{
    "0011/astro-ph0011134_arXiv.txt": {
        "abstract": "Ground-based \\& HST images of the nearby galaxy SBS~1543+593 ($z=0.009$) show it to be a Low Surface Brightness (LSB) galaxy with a central surface brightness of $\\mu_B(0)\\:=23.2$ mag arcsec$^{-2}$ and scale length 0.9~\\h , values typical for the local LSB galaxy population.  The galaxy lies directly in front of the QSO HS~1543+5921 ($z=0.807$); an HST STIS spectrum of the quasar reveals a damped \\lya\\ (D\\lya ) line at the redshift of the interloper with an H~I column density of $\\log N$(H~I)$\\:=\\:20.35$, as well as several low-ionization metal lines with strengths similar to those found in the Milky Way interstellar medium. Our data show that LSB galaxies are certainly able to produce the D\\lya\\ lines seen at higher redshift, and fuels the speculation that LSB galaxies are a major contributor to that population of absorbers. ",
        "introduction": "In the course of spectroscopic follow-up observations of quasar candidates in the Hamburg Quasar Survey, a redshift of $z=0.807$ was measured for the QSO HS~1543+5921.  However, Schmidt-plate images showed the object to be very extended, suggesting that the QSO might be centered on a low redshift galaxy.  Follow-up observations by \\citet[hereafter RH98]{RH98} found the QSO to be aligned with the foreground galaxy SBS~1543+593, and measurement of an H~II region in a spiral arm revealed the galaxy to be at $z=0.009$.  Although detecting QSOs close to nearby galaxies is not difficult, finding one which shines through the center of a galaxy and which is bright enough to be observed spectroscopically with HST is extremely rare. In this paper we present ground-based and HST optical images of SBS~1543+593 and show that the galaxy is actually a Low Surface Brightness (LSB) galaxy. We also present an HST spectrum of the background QSO HS~1543+5921, which reveals a damped \\lya\\ (D\\lya) absorption line at the redshift of the foreground galaxy. ",
        "conclusions": "Our optical and spectroscopic data have shown that SBS~1543+593 is a LSB galaxy causing a D\\lya\\ system at $z=0.009$ in the spectrum of HS~1543+5921. This makes it the lowest redshift D\\lya\\ system discovered outside of the local group. The identification of D\\lya\\ from a {\\it known} LSB is important, for the following reasons.  So far, results from ground-based and HST imaging of fields around QSOs known to show $z<1$ D\\lya\\ lines have been surprising, with the detection of a whole variety of galaxy types, including normal early and late-type HSB spirals, and amorphous LSB galaxies, identified as responsible for the absorption \\citep{steidel95, steidel94, Lanz97, LeBr97, RaTu98, pett00}.  This wide variety of absorber types has led to the speculation that D\\lya\\ systems may not simply signal the presence of normal gas-rich spiral galaxies after all, as has been postulated since their discovery \\citep{wolfe86}. It is important to note, however, that in most cases there exists no redshift information for the purported absorbers. Proximity to the line of sight is no guarantee that an `identified' object is the absorber, since there are often several absorption systems at redshifts other than that of the D\\lya\\ line along the QSO line of sight. If galaxies are responsible for these other systems, then the chance of mis-identification is high (particularly if D\\lya\\ systems do not arise in normal galaxies).  Hence the detection of a D\\lya\\ line from SBS~1543+593 is unique, in that  we know unequivocally that the absorber is an LSB galaxy.   It also seems likely that if SBS~1543+593 were moved to a redshift similar to those of the $z<1$ D\\lya\\ systems already studied, it would be extremely difficult to detect, partly due to its low surface brightness and partly due to its close proximity to the QSO and its small angular size. In fact, in this case, it is more likely that the nearby barred spiral galaxy south-west of the pair, which we take to be background to SBS~1543+593, would be identified as the D\\lya\\ absorber if no redshift information were available.    Our observations of SBS~1543+593 support the idea that LSB galaxies may contribute significantly to the population of D\\lya\\ absorbers, as initially suggested by \\citet{Imp89}. Although finding a D\\lya\\ from an LSB galaxy does not prove that {\\it all} D\\lya\\ systems are LSB galaxies, our detection does prove that LSB galaxies can produce such systems.  Perhaps more significant is the detection of this QSO-galaxy pair in the first place. Finding any bright QSO shining through the center of a nearby galaxy is extremely rare, and it is intriguing that an LSB galaxy is the interloper.  LSB galaxies are believed to be relatively free of dust, and hence optically thin to the photons of background objects --- these may therefore be the best types of galaxies to allow the light of quasars to shine through. This potential selection effect has been a concern for interpreting the copious abundance measurements of high-redshift D\\lya\\ systems: if a particular type of galaxy preferentially favors QSO light passing through it, then the derived metallicities are only applicable for that type of galaxy. There are also theoretical bases for the idea that LSB galaxies hold most of the high column density H~I cross-section: for example, \\citet{raul99} predicted that HSB disks consume neutral gas too fast to explain the observed evolution in the neutral gas mass density with redshift, and that consumption of hydrogen by LSB galaxies better fits the abundance measurements.  The mere existence of the QSO-galaxy pair discussed herein seems to strongly support these views. Unfortunately, just how common such QSO/LSB-galaxy alignments are at higher redshift will always be difficult to determine, due to the intrinsic faintness of the galaxies, especially if the galaxy is close to the QSO sightline as with this pair. Finally, we note that HS~1543+5921 is bright enough to be used to measure abundances in the interstellar medium of SBS~1543+593 with further HST observations. The current data do not have sufficient resolution for reliable metallicity estimates, so follow-up UV observations would be extremely valuable. Comparison of gas metallicities from a known LSB galaxy with higher redshift D\\lya\\ systems would be useful in deciding if LSB galaxies are, in general, responsible for the absorption. Measurement of the alpha-to-iron elemental abundances in SBS~1543+593 could also be used to derive the star-formation history of the galaxy, and relative metal abundances may provide insight into its dust content. Lastly, comparison of the kinematic structure of metal lines arising from SBS~1543+593 could be matched to global H~I kinematics derived from 21~cm emission maps of the galaxy, and to profiles seen in high-redshift D\\lya\\ lines which are believed by some authors to be indicative of rotating disks."
    },
    "0011/astro-ph0011072_arXiv.txt": {
        "abstract": "We present here numerical hydrodynamic simulations of line-driven accretion disk winds in cataclysmic variable systems. We calculate wind mass-loss rate, terminal velocities, and line profiles for \\ion{C}{4} (1550~\\AA) for various viewing angles. The models are 2.5-dimensional, include an energy balance condition, and calculate the radiation field as a function of position near an optically thick accretion disk. The model results show that centrifugal forces produce collisions of streamlines in the disk wind which in turn generate an enhanced density region, underlining the necessity of two dimensional calculations where these forces may be represented. For disk luminosity $L_{disk}=L_{\\sun}$, white dwarf mass $M_{wd}=0.6M_{\\sun}$, and white dwarf radii $R_{wd}=0.01R_{\\sun}$, we obtain a wind mass-loss rate of $\\dot M_{wind}=8 \\times 10^{-12} M_{\\sun} {\\rm yr}^{-1}$, and a terminal velocity of $\\sim 3000 {\\rm \\; km \\; s}^{-1}$. The line profiles we obtain are consistent with observations in their general form, in particular in the maximum absorption at roughly half the terminal velocity for the blue-shifted component, in the magnitudes of the wind velocities implied by the absorption components, in the FWHM of the emission components, and in the strong dependence in inclination angle. ",
        "introduction": "The first evidence for winds from cataclysmic variables (CVs) came from the discovery of P~Cygni profiles in the UV resonance lines of SS~Cygni (\\cite{hea85}). The dependence on the observability of CV winds with inclination angle and the apparent similarity between cataclysmic variable line profiles and those of OB stars led to the early suggestion by C\\'ordova \\& Mason (1982) that the winds in CVs originate from the accretion disk and that the line radiation pressure is responsible for the wind. More recently, P~Cygni profiles have been detected from virtually all nonmagnetic CVs with high mass accretion rates ($\\gtrsim 4\\times 10^{-10} M_{\\sun} {\\rm yr}^{-1}$) (e.g., \\cite{cor82}; \\cite{gre82}; \\cite{gui82}; \\cite{pri95}; \\cite{fri97}; \\cite{gan97}; \\cite{kni97}). The absorption component is most apparent in low inclination systems, and is not detected in eclipsing systems (e.g., \\cite{hut80}; \\cite{kra81}; \\cite{hol82}; \\cite{cor85}; \\cite{mas95}; \\cite{che97}).  Past efforts in the development of models for CV winds focused on one-dimensional models (\\cite{vit88}; \\cite{kal88}), and on kinematic modeling (\\cite{shl93}; \\cite{kni95}), which succeeded in showing consistency between the assumed polar geometry of a disk wind and observed profiles.  Icke~(1980) developed a two-dimensional disk wind model but did not take into account the radiation pressure due to line absorption. For a typical white dwarf mass of $0.6M_{\\sun}$ (\\cite{lei80}) the disk luminosity required to produce such a wind without line radiation pressure (assuming radiation pressure due to continuum scattering only) would require a luminosity of $\\sim 10^5 L_{\\sun}$ (\\cite{ick80}), several magnitudes above the observed luminosity of CVs which vary between $0.01L_{\\sun}$ and $10 L_{\\sun}$ (\\cite{pat84}). Icke did find that for accretion disk winds in general, for sufficiently high disk luminosity, a biconical disk wind would be produced. Icke~(1981) suggested that biconical winds were a general property of accretion disk winds independent of the wind driving mechanism.  Murray et~al. (1995) found that for line-driven disk winds in active galactic nuclei (AGN) the wind flow tends to be parallel to the accretion disk and developed a one-dimensional disk wind model for these systems. Murray \\& Chiang (1996) found that the single-peaked optical emission lines seen in AGNs and high-luminosity CVs could be accounted for by emission at the base of an accretion disk wind. Using the disk wind model of Murray et~al. (1995), Murray \\& Chiang (1997) calculated synthetic line profiles of \\ion{C}{4} 1550\\AA \\ and found qualitative agreement between their models and observed line profiles. In the case of CVs, as we discussed above, the winds are most apparent in low-inclination systems and not observed in eclipsing systems, indicating that the winds in CVs tend to flow perpendicular to the accretion disk.  In an earlier paper (\\cite{per97}, hereafter Paper~I) we presented two-dimensional isothermal hydrodynamic models of line-driven accretion disk winds (LDADW) in CVs. To our knowledge these isothermal models were the first two-dimensional hydrodynamic models developed for LDADWs. Results from Paper~I show, in analogy with line-driven winds from early type stars, that terminal velocities are approximately independent of the luminosity of the disk, although increments in luminosity produce increments in mass-loss rate. In Paper~I we showed that rotational forces are important in the study of winds from accretion disks, and that they cause the velocity streamlines to collide. The collision of streamlines reduce the speed and increase the density of the wind producing an enhanced density region. In Paper~I we also showed that the highest absorption occurs in the enhanced density region where density is increased relative to a spherically diverging wind with the same mass loss rate and the velocity is roughly half the terminal velocity. This density increase is necessary in order to produce at least marginally optically thick lines.  Recently Proga, Stone, \\& Drew (1998) also developed two-dimensional models for these systems, obtaining similar results to the models presented here. When evaluating the line-radiation pressure from the disk, in order to make the computations manageable, they assumed that the velocity gradient is primarily along the ``z'' direction, as we also did in Paper~1 and in the models presented here. Proga, Stone, \\& Drew worked out their hydrodynamic model in spherical coordinates, while the models presented here were implemented in cylindrical coordinates. They found similar wind velocities and wind mass loss rates as we do here, although they assumed disk luminosities about an order of magnitude greater. A difference we find with the results of the models presented by Proga, Stone \\& Drew (1998) is that they obtain unsteady flows characterized by large amplitude fluctuations in velocity and density, while for the models presented in this paper we find a steady flow.  A shortcoming of our previous work, presented in Paper~I, was that we had assumed an isothermal wind and an isothermal disk. In this work we develop a hydrodynamic model of LDADWs in CVs which includes the radial structure of an optically thick accretion disk with the corresponding radiation fields and surface temperature distributions. We have also implemented the corresponding energy conservation equation for the models presented here self consistently (see equation~[18]), including the adiabatic heating and cooling effects due to compression and expansion.  For the models presented in this paper we have also taken values presented by Abbott~(1982) for the line radiation force multiplier parameters, rather than the earlier values used by Castor, Abbott, \\& Klein~(1975) as we did in our isothermal models presented in Paper~I.  From the results of the hydrodynamic model presented in this paper, we calculate theoretical line profiles for the \\ion{C}{4} 1550~\\AA \\ line, and find that the line profiles obtained through our model are consistent with observations in their general form and strong dependence with inclination angle. Thus we are able to predict the general observed wind properties of CVs.  The boundary layer where the disk intercepts the white dwarf and the white dwarf photosphere may also contribute to the radiation field. The white dwarf by itself (without considering the effects of the accreting mass) will have a luminosity of the order of $0.01L_{\\sun}$(\\cite{lei80}), while the accretion disk will have typical luminosities of the order of $L_{\\sun}$. In a steady state accretion disk half of the energy of the accreting mass is converting into radiation emitted by the accretion disk. Assuming a relatively slowly rotating compact star, it has been predicted that the boundary layer presents a luminosity approximately equal to the accretion disk, thus accounting for the other half of the energy of the accreting mass (e.g., \\cite{fra92}). But there is still uncertainty concerning the existence and spectrum of the boundary layer (\\cite{vrt94}; \\cite{mau95}; \\cite{lon96}). Analysis of observations have found, in some cases, a boundary layer luminosity an order of magnitude less than the disk luminosity (e.g., \\cite{mau95}), and in other cases a boundary layer luminosity comparable to disk luminosity (\\cite{lon96}). In the models presented here we have assumed that the radiation field is generated by the accretion disk alone and have therefore neglected any boundary layer or white dwarf radiation. In the future we plan to explore the effects of boundary layer radiation on our models.  In \\S~2 we discuss the radial structure of the accretion disk and the radiation field as implemented in our model. In \\S~3 we derive the expression we use for the treatment of the line radiation pressure. In \\S~4 we present and discuss the hydrodynamic calculations. In \\S~5 we present and discuss theoretical \\ion{C}{4} 1550~\\AA \\ line profiles and compare them with observations. In \\S~6 we present a summary and conclusions of this work. ",
        "conclusions": "We have developed a 2.5-dimensional hydrodynamic line-driven adiabatic accretion disk wind model. Our model solves a complete set of adiabatic hydrodynamic partial differential equations, using the PPM numerical scheme and implementing the radial temperature and radiation emission distributions on the surface of an accretion disk. From our models we calculate wind mass-loss rates, terminal velocities, and line profiles for \\ion{C}{4} (1550~\\AA) for various angles. For typical cataclysmic variable parameter values of disk luminosity $L_{disk}=L_{\\sun}$, white dwarf mass $M_{wd}=0.6M_{\\sun}$, and white dwarf radii $R_{wd}=0.01R_{\\sun}$, we obtain a wind mass-loss rate of $\\dot M_{wind}=8 \\times 10^{-12} M_{\\sun} {\\rm yr}^{-1}$, and a terminal velocity of $\\sim 3000 {\\rm \\;km \\;s}^{-1}$. Studies of observed ultraviolet line profiles from CVs have estimated CV  wind mass loss rates to be in the order of $10^{-12}$ to $10^{-11} M_{\\sun} {\\rm yr}^{-1}$ and found CV wind terminal speeds in the order of $\\sim 3000 \\; {\\rm km} \\; {\\rm s}^{-1}$ (\\cite{kra81}; \\cite{cor82}; \\cite{gre82}; \\cite{mau87}; \\cite{pri95}; \\cite{fri97}), thus the models developed in this work predict values for wind mass loss rates and terminal speed consistent with observations.  Furthermore, the line profiles we obtain through our computational model are also consistent with observations in their general form, in particular in the maximum absorption at roughly half the terminal velocity for the blue-shifted component, in the magnitudes of the wind velocities implied by the absorption components, in the FWHM of the emission components, and in the strong dependence in inclination angle.  We have shown that rotational forces are important in the study of winds from accretion disks. They cause the velocity streamlines to collide which results in an enhanced density region. In this region the wind speed is reduced and the wind density is increased. The increase in density caused by the collision of streamlines is important because it permits the appearance of blue-shifted absorption lines as observed in P~Cygni profiles of low-inclination CVs.  A shortcoming of our models is the neglect of radiative cooling and heating and local ionization equilibrium throughout the wind. In the next paper of this series we will present models which include these effects. We also plan to extend our models to the study of the wind dynamics of low mass X-ray binaries and active galactic nuclei, where these processes may play a more important role in determining the overall dynamics."
    },
    "0011/astro-ph0011244_arXiv.txt": {
        "abstract": "{ We present the results of a study based on an \\XMM\\ Performance Verification observation of the central 30\\arcmin\\ of the nearby spiral galaxy M31.  In the 34-ks European Photon Imaging Camera (EPIC) exposure, we detect 116 sources down to a limiting luminosity of \\scinot{6}{35}~\\Lcgs\\ (0.3--12~keV, $d=760$~kpc).  The luminosity distribution of the sources detected with \\XMM\\ flattens at luminosities below $\\sim2.5 \\times 10^{37}$~\\Lcgs.  We make use of hardness ratios for the detected sources in order to distinguish between classes of objects such as super-soft sources and intrinsically hard or highly absorbed sources. We demonstrate that the spectrum of the unresolved emission in the bulge of M31 contains a soft excess which can be fitted with a $\\sim$0.35-keV optically-thin thermal-plasma component clearly distinct from the composite point-source spectrum.  We suggest that this may represent diffuse gas in the centre of M31, and we illustrate its extent in a wavelet-deconvolved image. ",
        "introduction": "\\begin{figure*} \\vbox{\\psfig{figure=fig1half.ps,width=12cm}\\vspace{-8.0cm}} \\hfill\\parbox[b]{5.5cm}{ \\caption{Three-colour \\XMM\\ EPIC MOS1 image of the central 30\\arcmin$\\times$30\\arcmin\\ of M31.  The red, green, and blue intensities correspond to logarithmically-scaled counts from energy bands of 0.3--0.8~keV, 0.8--2~keV, and 2--8~keV respectively.  The image was constructed with 2\\arcsec\\ pixels and has been smoothed with a Gaussian of FWHM 4.4\\arcsec, approximately equal to that of the MOS1 PSF. \\vspace{4.0cm} } } \\label{fig:image} \\end{figure*}   Being the closest spiral galaxy to our own, the Andromeda Galaxy (M31) is in many respects ideal for the study of X-ray emission in a galaxy similar to the Milky Way.  The sources in M31 are observed at a nearly uniform distance, and, owing to the inclination of the galaxy (77$^{\\circ}$), they are viewed through a substantially lower absorption column than for sources in the Galactic plane.  In a recent review, van den Bergh (2000) reports a distance modulus to M31 of 24.4$\\pm$0.1, corresponding to a distance of 760~kpc. We adopt this value in our analysis, and for consistency we scale to this distance when discussing published luminosities which assume a different distance.  Over 100 X-ray sources in M31 were detected with the \\Einstein\\ observatory (Trinchieri \\& Fabbiano 1991; van Speybroeck et al.\\ 1979).  The brightest X-ray source in M31 was found to have a luminosity of $\\sim\\mscinot{3}{38}$~\\Lcgs, approximately the Eddington luminosity for spherical accretion onto a 1.4~\\msun\\ neutron star.  For sources down to $\\mscinot{2}{36}$~\\Lcgs, the luminosity distribution was reported to be consistent with a single power law which, extrapolated to fainter levels, could fully account for the X-ray emission from the bulge of M31.  Primini, Forman \\& Jones (1993) detected 86 X-ray sources in the central 34\\arcmin\\ of M31 with the \\ROSAT\\ HRI.  They found a break in the luminosity distribution at $\\sim\\mscinot{2}{37}$~\\Lcgs, below which the distribution of sources flattened.  This flattening suggested that the detected population of X-ray sources could account for only $\\sim$15--26\\% of the unresolved X-ray emission in M31. Contributions from known less-luminous populations of X-ray sources also could not fully account for the unresolved emission, suggesting that the remaining emission is truly diffuse or due to a new class of X-ray sources.  In an extensive, 6.3 deg$^2$, survey of M31 with the \\ROSAT\\ PSPC, 396 X-ray sources were detected (Supper et al.\\ 1997).  However, only 22 of these sources were detected in the bulge region ($r < 5$\\arcmin) due to the resolution of the PSPC.  In the first \\Chandra\\ observation of M31, the nuclear source seen with \\Einstein\\ and \\ROSAT\\ was resolved into five sources (Garcia et al.\\ 2000).  One of these sources is located within 1\\arcsec\\ of the radio nucleus of M31 and exhibits an unusually soft X-ray spectrum, suggesting that it may be associated with the central super-massive black hole.  A few more pairs of previously unresolved sources and a new transient were also detected within 30\\arcsec\\ of the nucleus.  M31 was selected as an \\XMM\\ (Jansen et al.\\ 2001) Performance Verification target in order to demonstrate the capabilities of \\xmm\\ in performing spectral and timing studies in a field of point sources and extended emission. In this paper we focus on the group properties of the X-ray point sources in M31 as well as the diffuse emission.  In a companion paper (Osborne et al.\\ 20001, Paper II), we discuss the spectral and timing properties of individual X-ray sources in M31. ",
        "conclusions": "Using \\XMM\\ data, we have confirmed that the X-ray emission from the bulge of M31 is dominated by bright point-like sources, most of which are likely to be low-mass X-ray binaries.  For sources in the central region of M31, we have confirmed that the luminosity distribution is flatter toward lower luminosities (Primini et al.\\ 1993).  The steepening of the luminosity distribution above $2.5 \\times 10^{37}$\\Lcgs\\ is indicative of a lack of bright sources in M31 (cf.\\ the source distribution in M33, Long et al.\\ 1996). Only two sources in our sample have a $0.3$--$12$ keV unabsorbed luminosity $\\geq 10^{38}$\\Lcgs.  As in previous observations with \\Einstein\\ and \\ROSAT\\ (Trinchieri \\& Fabbiano 1991; Primini et al.\\ 1993), significant unresolved emission was found to contribute to the total emission of the bulge. The flattening of the luminosity distribution for fainter sources means that an extrapolation of the detected population of point sources at lower energies cannot account for the total core emission of M31.  A soft excess in the spectrum of the M31 bulge was previously reported in \\ROSAT\\ and \\BeppoSAX\\ observations (Irwin \\& Bregman 1999; Trinchieri et al.\\ 1999).  Our analysis of the \\XMM\\ data shows that the soft component in the spectrum of the bulge is associated with unresolved emission; this confirms the results of Borozdin \\& Priedhorsky (2000) based on \\ROSAT\\ data.  More importantly, our \\XMM\\ study has revealed for the first time that while the integrated spectrum of point-like sources is featureless, the spectrum of the unresolved emission shows multiple emission lines typically found in the spectrum of hot, optically thin plasma. Therefore, we suggest that the second significant source of X-ray emission in the bulge is truly diffuse gas with an effective temperature $\\sim$0.35 keV.  The contribution of this gas to the total unabsorbed X-ray luminosity is estimated to be $\\sim$10\\% in 0.3--12~keV band (corresponding to $\\approx 2 \\times 10^{38}$\\Lcgs), but more than 20\\% in the \\ROSAT\\ band (0.1--2.4~keV).  The significance of this result goes far beyond the case of M31, because the bulge of this galaxy is often considered as a prototype for the population of early-type X-ray galaxies. For example, Sarazin et al.~(2000) recently reported that, according to \\Chandra\\ observations, 23\\% of X-ray emission from NGC\\,4697 is emitted by interstellar gas, contrary to their previous expectations (Irwin et al.\\ 2000).  Two more \\XMM\\ observations of the central region of M31 are scheduled as part of the Guaranteed-Time program, as are observations of five additional fields along the disk of M31.  These will allow us to reach fainter flux levels in the bulge and to study the populations of X-ray sources in different parts of the M31 galaxy."
    },
    "0011/hep-ph0011353.txt": {
        "abstract": "We review new results in neutrino physics, including the latest data of the Super-Kamiokande, SNO and K2K experiments. ",
        "introduction": "The recent evidence for neutrino oscillations \\cite{SK1} provided us with the first firm evidence of physics beyond the standard model and opened a new an exciting era in neutrino studies. Neutrino physics is a very active branch of particle physics now, both experimentally and theoretically. The experimental data keep pouring in, and new experiments are either under way or in an advanced stage of planning. On the theoretical side, there are new analyses of the data, both in the 3-flavour and 4-flavour schemes; the analyses of the solar neutrino data are being extended to cover the ``dark side'' of the parameter space, not studied (or little studied) before; the reach of future planned experiments is being investigated and new experiments designed to eliminate the white spots on the map of neutrino properties are being suggested. In addition, since the discovery of neutrino oscillations a large number of models of neutrino mass was proposed and some old models were reconsidered in the light of the new data.  In the present lectures we review new results in neutrino physics, mainly the experimental ones. We discuss the latest results of the solar and atmospheric neutrino experiments and their analyses; results from the reactor and accelerator experiments, including those from the first long-baseline accelerator neutrino experiment K2K, and then briefly discuss future experiments and projects. Finally, we discuss how all the presently available data can be summarized concisely in terms of the phenomenologically allowed structures of the neutrino mass matrix.  We do not discuss neutrino mass model because of the lack of space; for recent reviews on this subject the reader is referred to \\cite{Sm1,AltFer,Tan,Moh1,Fr,Barr1}.  The data reviewed in the lectures given at the NATO ASI in Cascais included the results reported at the Neutrino 2000 Conference in Sudbury, June 16 - 21, 2000. However, the present written version is updated to include the results reported up to November of 2000. %up to and including those reported at the ICHEP in Osaka, July 27 - August %2, 2000. For other recent reviews on neutrino physics see, e.g., \\cite{Bil,Akh1}. ",
        "conclusions": ""
    },
    "0011/astro-ph0011522_arXiv.txt": {
        "abstract": "We combine our compilation of photometry of M31 globular clusters and probable cluster candidates with new near-infrared photometry for 30 objects. Using these data we determine the globular cluster luminosity function (GCLF) in multiple filters for the M31 halo clusters. We find a GCLF peak and dispersion $V_0^0=16.84\\pm0.11$, ${\\sigma}_t=0.93\\pm0.13$ (Gaussian ${\\sigma}=1.20\\pm0.14$), consistent with previous results. The halo GCLF peak colors (e.g. $B_0^0-V_0^0$) are consistent with the average cluster colors. We also measure $V$-band GCLF parameters for several other subsamples of the M31 globular cluster population. The inner third of the clusters have a GCLF peak significantly brighter than that of the outer clusters ($\\Delta V^0\\approx 0.5$~mag). Dividing the sample by both galactocentric distance and metallicity, we find that the GCLF also varies with metallicity, as the metal-poor clusters are on average 0.36~mag fainter than the metal-rich clusters. Our modeling of the catalog selection effects suggests that they are not the cause of the measured GCLF differences, but a more-complete, less-contaminated M31 cluster catalog is required for confirmation. Our results imply that dynamical destruction is not the only factor causing variation in the M31 GCLF: metallicity, age, and cluster initial mass function  may also be important. ",
        "introduction": "Globular clusters (GCs) are unique markers of the formation and evolution of galaxies. They are bright, easily identifiable packages of Population~II stars with homogeneous abundances and history. GCs are relics of the earliest star formation in galaxies and have witnessed many changes in their parent galaxies. The study of globular cluster systems (GCSs) is thus important in understanding the history of galaxies.  The distribution of integrated GC magnitudes, known as the globular cluster luminosity function (GCLF), has long been known to be unimodal and approximately symmetric in the Milky Way. The assumption that these properties are universal has allowed the determination of GCLF parameters for over a hundred other galaxies, and the peak absolute magnitude is found to be roughly constant from galaxy to galaxy \\citep{har88}. Since the variation in mass-to-light ratios among GCs is fairly small \\citep[see, e.g.,][]{dg97} the constant peak magnitude implies the existence of a characteristic mass scale. Theorists have attempted to explain why there should be a characteristic mass scale for globular clusters: is it a property of formation \\citep[e.g.][]{pd68}, a result of subsequent dynamical processes \\citep[e.g.][]{ot95}, or a combination? The constant peak magnitude also presents a challenge to observers, who have attempted to quantify the variation in peak magnitude by environment \\citep{bt96}, galaxy type and luminosity \\citep{whi97}, and color \\citep{kws99}. Many observers have also attempted to use the GCLF peak as a standard candle for distance measurement. This method has had a mixed reception, with some authors \\citep[e.g.][]{whi97} claiming good results and others \\citep[e.g.][]{fer99} less complimentary.  GCLF measurements are made in a number of different bandpasses. The most widely used are $V$ and $B$, (and their Hubble Space Telescope (HST) equivalents, F555W and F450W); the I-band equivalent F814W is also commonly used for HST data. $B$ and $V$ band mass-to-light ratios are sensitive to metallicity, so GCSs with the same underlying mass distribution but different average metallicities will have different GCLFs. \\citet*{acz95} used the relation between metallicity and mass-to-light ratio as predicted by \\citeauthor{w94}'s (1994) simple stellar population models to estimate the effects of metallicity variations on the GCLF peak. Metal-rich clusters should be fainter (their $M/L$ is larger), and the effect on the GCLF can be substantial. For example, \\citet{acz95} predict that a change in mean GCS metallicity from $\\overline{\\rm [Fe/H]}=-1.35$~dex (the Milky Way value) to $-0.60$~dex \\citep[the value for the elliptical NGC~3923;][]{z95} shifts the GCLF peak by $\\Delta B^0=0.35$ and $\\Delta V^0=0.22$. The same metallicity change shifts the $J$ band GCLF peak by $\\Delta J^0=-0.09$, because the metallicity effect on $M/L$ changes direction for bandpasses redward of $I$. Observing at longer wavelengths also reduces the effects of extinction, both Galactic and within the GCS parent galaxy.  Destruction of globular clusters through dynamical effects could produce a globular cluster mass function (GCMF) that varies with distance from the center of the host galaxy. The three effects usually considered by modelers are disk shocking, evaporation, and dynamical friction; the first two destroy low-mass, low-concentration clusters, while the last is most effective for massive clusters. All three effects are strongest at small Galactocentric distance $R_{gc}$. Even at small $R_{gc}$, the dynamical friction timescale for typical-mass GCs is much longer than a Hubble time, and dynamical friction is probably not important in GCMF evolution \\citep{ot95}. Low-mass clusters near the Galactic center are thus most prone to destruction. Assuming a mass-to-light ratio independent of $R_{gc}$ -- if the GCS has a radial metallicity gradient this is not strictly true since $M/L$ depends on [Fe/H] -- the radial difference in the GCMF can be translated into a radial difference in the GCLF. Many authors \\citep{b97,og97,v98} have predicted the size of the GCLF variation, with widely differing results (see Section~\\ref{sec-implic}).   In \\citet{b00} we compiled the best available photometry for 435 M31 globular clusters and plausible candidates. This catalog contains $V$ magnitudes for almost all objects and $B$ magnitudes for about 90\\%, but completeness in the longer-wavelength bands is much lower. About 55\\% of the objects have $I$ magnitudes, and the same fraction, although not the same objects, have $JK$; Figure~\\ref{magcomp-ir} shows the completeness of the existing IR photometry. In this paper we estimate the parameters of the halo clusters' GCLF in six different bandpasses; this includes the first GCLF measurements in the near-IR. We then investigate the variation in the $V$-band GCLF parameters for several different subsamples of M31 GCs. Finally, we consider the implications of the measured GCLF variations for GCS and galaxy formation and evolution. ",
        "conclusions": "We have calculated the first $URJK$ GCLFs for M31 halo globular clusters, and find that the GCLF peak colors are consistent with the average cluster colors. Our parameters for the $V$- and $B$-band halo GCLFs are consistent with those of other groups. We find no significant differences between the disk and halo GCLF peaks, although the disk has a lower dispersion. A difference in GCLF peak at the $2\\sigma$ level occurs when we consider the inner and outer-most groups, as determined by projected galactocentric distance. This difference is consistent with that predicted by \\citet{og97} for M31; however, we do not detect the predicted difference in GCLF dispersion. We separate the M31 clusters by metallicity and find that the metal-rich clusters have a brighter GCLF peak than the metal-poor clusters, even when the difference in $R_{gc}$ is taken into account. Modeling of the catalog selection effects suggests that these effects are not responsible for the measured GCLF differences. However, an M31 GC catalog with well-understood and spatially uniform completeness and contamination is required in order to definitively confirm our results. Such a catalog might be produced by a near-IR, high spatial resolution survey of M31. We consider the implications of the GCLF differences for models of globular cluster and GCS formation, and conclude that younger ages for metal-rich clusters plus dynamical destruction of inner clusters are the most likely causes of the observed GCLF variations."
    },
    "0011/astro-ph0011008_arXiv.txt": {
        "abstract": "We present optical spectroscopic identification of sources identified in the BeppoSAX MECS fields of the High Energy LLarge Area Survey (HELLAS). In total 62 sources from a sample of 115 brighter than $F_{5-10keV}>5\\times10^{-14}$ $cgs$ have been identified. We find a density of 13-21 sources/deg$^2$ at $F_{5-10keV}>5\\times10^{-14}$ $cgs$, which contribute to $20-30$ \\% of the hard X-ray background. Evidences are found for type 1 AGN being more absorbed with increasing redshift or luminosity. The low redshift ($z<0.2$) ratio of type 2 to type 1 AGN is $3\\pm1.5$, in agreement with the unified models for AGN. The luminosity function of type 1 AGN in the 2-10 keV band is preliminary fitted by a two-power law function evolving according to a pure luminosity evolution model: $L\\propto (1+z)^{2.2}$. ",
        "introduction": "Hard X-ray observations are the most efficient way of tracing emission due to accretion mechanisms, such in Active Galactic Nuclei (AGN), and sensitive hard X-ray surveys are powerful tools to select large samples of AGN less biased against absorption and extinction.  In this framework we decided to take advantage of the large field of view and good sensitivity of the BeppoSAX MECS instrument (Boella et al. 1997a,b) to survey tens to hundreds of square degrees at fluxes $\\ga 5-10\\times10^{-14}~ erg~cm^{-2}s^{-1}$ (Fiore et al. 2000a), and using higher sensitivity XMM-Newton and Chandra observations to extend the survey down to $\\sim10^{-14}~ erg~cm^{-2}s^{-1}$ on several deg$^2$. The results from the optical identification of a sample of faint Chandra sources discovered over the first 0.14 deg$^2$ have been published by Fiore et al. (2000b). This approach is complementary to deep pencil beam surveys ($\\sim0.1$ deg$^2$, see e.g. Mushotzky et al. 2000, Hornschemeier et al. 2000), as we cover a different portion of the redshift--luminosity plane. Our purpose is to study cosmic source populations at fluxes where the X-ray flux is high enough to provide X-ray spectral information in higher sensitivity follow-up observations. This would allow the determination of the distribution of absorbing columns in the sources making the hard XRB, providing strong constraints on AGN synthesis models for the XRB (e.g. Comastri et al. 1995). ",
        "conclusions": ""
    },
    "0011/gr-qc0011109_arXiv.txt": {
        "abstract": "The development of both ground- and space-based gravitational wave detectors provides new opportunities to observe the radiation from binaries containing neutron stars and black holes.  Numerical simulations in 3-D are essential for calculating the coalescence waveforms, and comprise some of the most challenging problems in astrophysics today.  This article briefly reviews the current status of efforts to calculate black hole and neutron star coalescences, and highlights challenges for the future. ",
        "introduction": "Binaries containing black holes (BHs) or neutron stars (NSs) are among the most important and interesting sources for the gravitational wave detectors expected to start operation in the early part of the $21^{\\rm st}$ century.  These binaries spiral together due to the emission of gravitational radiation, leading to the final collision and coalescence of their components.  The NS/NS, NS/BH, and stellar BH/BH binaries are target sources for ground-based interferometers such as LIGO, VIRGO, GEO600, and TAMA, while the space-based LISA is expected to be most sensitive to massive BH binaries (see the articles by Weiss and Kalogera in these Proceedings). Astrophysically, the data from these detectors can provide important insights into NS properties; the equation of state of matter at nuclear densities; models of gamma-ray bursts; active galactic nuclei and quasars; and the strong-field regime of gravity, including unambiguous detection of BH formation (Schutz 1997; Thorne 1998).  Calculations of the gravitational waves from such binaries focus on three physical regimes.  During the {\\em inspiral phase}, the components are widely separated and can be treated as point particles, allowing analytic calculations using post-Newtonian (PN) expansions.  Templates constructed from the resulting waveforms are expected to form a key component of data analysis and source identification (Flanagan \\& Hughes 1998). The {\\em coalescence phase} begins when the compact objects are close enough to suffer tidal deformation, and continues as a combination of relativistic and hydrodynamic effects drives the stars together on more rapid timescales. Calculation of the waveforms produced during this stage requires 3-D numerical simulations that solve the full Einstein equations and, for a NS, general relativistic hydrodynamics; in some cases, the formation of a BH must also be modeled.  Finally, the {\\em ringdown phase} encompasses the late-time behavior of the merger product or remnant, radiating in its normal modes (Thorne 1998).  Numerical simulation of the coalescence phase using full general relativity is one of the most challenging problems in astrophysics today.  This article surveys the current status of efforts to calculate coalescence waveforms for NS/NS, NS/BH, and BH/BH binaries. ",
        "conclusions": ""
    },
    "0011/astro-ph0011187_arXiv.txt": {
        "abstract": "We use hydrodynamical $N$-body simulations to study the kinetic Sunyaev--Zel'dovich effect. We construct sets of maps, one square degree in size, in three different cosmological models. We confirm earlier calculations that on the scales studied the kinetic effect is much smaller than the thermal (except close to the thermal null point), with an rms dispersion smaller by about a factor five in the Rayleigh--Jeans region. We study the redshift dependence of the rms distortion and the pixel distribution at the present epoch. We compute the angular power spectra of the maps, including their redshift dependence, and compare them with the thermal Sunyaev--Zel'dovich effect and with the expected cosmic microwave background anisotropy spectrum as well as with determinations by other authors. We correlate the kinetic effect with the thermal effect both pixel-by-pixel and for identified thermal sources in the maps to assess the extent to which the kinetic effect is enhanced in locations of strong thermal signal. ",
        "introduction": "The Sunyaev--Zel'dovich (SZ) effect (Sunyaev \\& Zel'dovich 1972, 1980; for reviews see Rephaeli 1995 and Birkinshaw 1999) is the change in energy experienced by cosmic microwave background photons when they scatter from intervening gas, especially that in galaxy clusters. The dominant version from clusters is the thermal SZ effect, the gain in energy acquired from the thermal motion of the gas which is commonly at a temperature of tens of millions of degrees in clusters. The kinetic SZ effect is the Doppler shift arising from the bulk motion of the gas.  The thermal effect has been quite well studied theoretically, and recently has become a burgeoning area of observational activity with the construction of two-dimensional maps of clusters becoming commonplace (Jones et al.~1993; Myers et al.~1997; Carlstrom et al.~2000).  By contrast, less attention has been given to the kinetic effect.  It presents a considerable observational challenge, because it is predicted to be much smaller than the thermal effect on the angular scales explored so far, and also because unlike the thermal effect it possesses no characteristic spectral signature allowing it to be distinguished from primary cosmic microwave background anisotropies.  It is also more difficult to make semi-analytic calculations. The thermal effect can be obtained fairly directly via the Press--Schechter (1974) approach; various aspects including number counts, the global CMB distortion, and its impact as secondary anisotropy for CMB measurements at small scales have been extensively studied (Cole \\& Kaiser~1988; Colafrancesco et al.~1994, 1997; Bartlett \\& Silk~1994; Barbosa et al.~1996; Eke, Cole \\& Frenk 1996; Aghanim et al.~1997; Komatsu \\& Kitayama 1999; Atrio-Barandela \\& M\\\"ucket 1999; Molnar \\& Birkinshaw 2000; Hern\\'andez-Monteagudo, Atrio-Barandela \\& M\\\"ucket 2000). Predictions for the kinetic effect require the simultaneous estimation of both mass and peculiar velocity, and the signal is much less dominated by the gas which happens to be in massive halos.  A detailed analytical model has recently been constructed by Valageas, Balbi \\& Silk (2000) [see also Benson et al.~2000], which computes the effect including inhomogeneities in reionization; this supersedes earlier modelling by Aghanim et al.~(1997) where the peculiar velocities were drawn from a gaussian distribution of fixed width.  We recently used large-scale hydrodynamical $N$-body simulations to make simulated maps of the thermal effect and analyzed their properties (da Silva et al.~2000; for similar recent work see Refregier et al.~2000a; Seljak, Burwell \\& Pen 2000; Springel, White \\& Hernquist~2000).  In this paper, we make a detailed analysis of maps of the kinetic effect, made using the same technique.  Recently Springel et al.~(2000) used similar simulations to study the angular power spectrum of the kinetic effect in the $\\Lambda$CDM cosmology, and investigated the influence of non-gravitational heating upon it. In this paper we study three different cosmologies, and investigate a wide range of properties of the resulting maps including the relation between thermal and kinetic distortion. ",
        "conclusions": "We have used hydrodynamical simulations to study a number of aspects of the kinetic SZ effect, including its dependence on cosmological parameters. We have studied the redshift dependence, pixel histograms, angular power spectra, and the correlation between the kinetic and thermal effects. We have confirmed that the kinetic effect has a dispersion more than a factor five below the thermal (in the Rayleigh--Jeans region), leading to an angular power spectrum a factor of typically twenty-five lower in the low-density cosmologies. For critical density we have found a smaller difference, but with much greater statistical uncertainty. For the thermal effect the angular power spectrum is mostly generated at redshifts below one, while for the kinetic effect a significant amount of the power has origin above redshift two. The correlation of the kinetic effect with the thermal confirms and quantifies the expected enhancement of the kinetic effect in regions with a strong thermal signal.  Our simulations are of ideal size to study the SZ effect on scales between one and several arcminutes, which is currently a resolution attracting great interest. Our results complement perfectly the small-scale simulation work of Bruscoli et al.~(1999) and Gnedin \\& Jaffe (2000), and recent semi-analytic work including that of Benson et al.~(2000) and Valageas et al.~(2000). The theory of these secondary anisotropies is now becoming highly developed; the key challenges for the kinetic effect lie very much on the observational side."
    },
    "0011/astro-ph0011378_arXiv.txt": {
        "abstract": "We show how estimates of parameters characterizing inflation-based theories of  structure formation localized over the past year when large scale structure (LSS) information from galaxy and cluster surveys was combined with the rapidly developing cosmic microwave background (CMB) data,  especially from  the recent Boomerang and Maxima balloon experiments. All current CMB data plus a relatively weak  prior probability on the Hubble constant, age and LSS points to little mean curvature ($\\Omega_{tot} = 1.08\\pm 0.06$) and nearly scale invariant initial fluctuations ($n_s =1.03 \\pm 0.08$), both predictions of (non-baroque) inflation theory. We emphasize the role that degeneracy among parameters in the $L_{pk} = 212\\pm 7$  position of the (first acoustic) peak plays in defining the $\\Omega_{tot}$ range upon marginalization over other variables. Though the CDM density is in the expected range ($\\Omega_{cdm}{\\rm h}^2=0.17 \\pm 0.02$), the baryon density $ \\Omega_b {\\rm h}^2 = 0.030\\pm 0.005$ is somewhat above the independent $0.019\\pm 0.002$ nucleosynthesis estimates. CMB+LSS gives independent evidence for dark energy ($\\Omega_\\Lambda = 0.66 \\pm 0.06$) at the same level as from supernova (SN1) observations, with a phenomenological quintessence equation of state limited by SN1+CMB+LSS to $w_Q < -0.7$ cf. the $w_Q$=$-1$ cosmological constant case. ",
        "introduction": "\\noindent {\\bf Experiments and Bandpowers:} Anisotropies at the $30 \\mu K$ level at low multipoles  revealed by COBE in 1992 were augmented at higher $\\ell$ in some 19 other experiments, some with a comparable number of resolution elements to the 600 or so for COBE, most with many fewer. A list of these experiments to April 1999 with associated bandpowers is given in Bond, Jaffe and Knox (2000 [BJK00]). The anisotropy picture dramatically improved this past year, as results were announced first in summer 99 from the ground-based TOCO experiment in Chile (Miller \\et 2000), then in November 99 from Boomerang-NA, the North American test flight (Mauskopf et 1999). These two additions improved peak localization and gave evidence for $\\Omega_{tot}\\sim 1$. Then in April 2000, results  from the first CMB long duration balloon (LDB) flight, were announced (de Bernardis \\et 2000), followed in May 2000 by results from the night flight of Maxima (Hanany \\et 2000). Boomerang's best resolution was $10^\\prime$, about 40 times better than that of COBE, with tens of thousands of resolution elements. Maxima had a similar resolution but covered an order of magnitude less sky. Fig.~1 shows the 150A GHz Boomerang-LDB map and the Wiener-filtered Maxima-1, to scale. The de Bernardis \\et (2000) maps at 90 and 220 GHz show the same spatial features as this  150 GHz one, with the overall intensities falling precisely on the CMB blackbody curve.  The Toco, Boomerang and Maxima experiments are described elsewhere in these proceedings. They were designed to reveal the {\\it primary} anisotropies of the CMB, those which can be calculated using linear perturbation theory. Fig.~1 shows the temperature power spectra for Boomerang,  Maxima and prior-CMB data (Boomerang-NA+TOCO+April 99) are in good agreement. Sketching the impact of these new results on cosmic parameter estimation (Lange \\et 2000 [Let00], Jaffe \\et 2000 [Jet00]) is the goal of this paper. Space constraints preclude adequate referencing here, but these are given in the Boomerang (Let00) and Maxima+Boomerang (Jet00) parameter estimation papers (see also Bond 1996, [B96], for other references).  \\begin{figure} \\vspace{-15pt} \\centerline{\\hspace{30pt}\\epsfig{file=figCLalldata.ps,height=4.5in}} \\vspace{-7pt} \\centerline{\\epsfig{file=figmap.ps,height=3.3in}} \\vspace{-12pt} \\caption{\\small The top figure shows ${\\cal C}_\\ell$ grouped in bandpowers for Boomerang-LDB (crosses), Maxima-I (triangles) and prior-CMB experiments (TOCO+Boomerang-NA+\"April 99\", squares). The lower panel contrasts the optimally-combined power spectra for Boomerang+Maxima+DMR (squares) with that for Boomerang+Maxima+prior-CMB (circles), showing the prior experiments do not move ${\\cal C}_\\ell$ very much. Best-fit models for arbitrary $\\Omega_{tot}$ and for $\\Omega_{tot}$=1 are shown in both panels. The Boomerang 150A GHz map (\\ie for one of 16 bolometers) and the multifrequency Wiener-filtered Maxima-I map, its 124 square degrees drawn to scale, are shown in the bottom figure. Only the 440 square degrees within the central rectangle of the entire 1800 square degrees covered by Boomerang were used in the analysis. \\normalsize } \\label{fig:CLdatmap} \\end{figure}  We are only at the beginning of the high precision  CMB era for primary anisotropies  heralded by the arrival of Boomerang and Maxima, with interferometers taking data (VSA, CBI, DASI), the single dish ACBAR about to, and new LDBs to fly in the next few years (Arkeops, Tophat, Beast/Boost), as well as Boomerang-2001 and the neo-Maxima Maxipol, both concentrating on polarization. In April 2001, NASA's HEMT-based MAP satellite will launch, with $12^\\prime$ resolution, and in 2007, ESA's bolometer+HEMT-based Planck satellite is scheduled for launch, with $5^\\prime$ resolution.  \\vskip 10pt \\noindent {\\bf The CMB Analysis Pipeline:} Analyzing Boomerang and other experiments involves a pipeline that takes (1) the timestream in each of the bolometer channels coming from the balloon plus information on where it is pointing and turns it into (2) spatial maps for each frequency characterized by average temperature fluctuation  values in each pixel (Fig.~1) and a pixel-pixel correlation matrix characterizing the noise, from which various statistical quantities are derived, in particular (3) the temperature power spectrum as a function of multipole (Fig.~1), grouped into bands, and two band-band error matrices which together determine the full likelihood distribution of the bandpowers (Bond, Jaffe \\& Knox 1998 [BJK98], BJK00). Fundamental to the first step is the extraction of the sky signal from the noise, using the only information we have, the pointing matrix mapping a bit in time onto a pixel position on the sky. To compare the data with millions of cosmological models, as we wish to do here, the radical compression step from 2 to 3 is essential, and hinges upon an accurate representation of the likelihood surface.  There is generally another step in between (2) and (3), namely separating the multifrequency spatial maps into the physical components on the sky: the primary CMB, the thermal and kinematic Sunyaev-Zeldovich effects, the dust, synchrotron and bremsstrahlung Galactic signals, the extragalactic radio and submillimetre sources. The strong agreement among the Boomerang maps indicates that to first order we can ignore this step, but it has to be taken into account as the precision increases.  The Fig.~1 map is consistent with a Gaussian distribution, thus fully characterized by just the power spectrum. Higher order (concentration) statistics (3,4-point functions, \\etc) tell us of non-Gaussian aspects, necessarily expected from the Galactic foreground and extragalactic source signals, but possible even in the early Universe fluctuations. For example, though non-Gaussianity occurs only in the more baroque inflation models of quantum noise, it is a necessary outcome of defect-driven models of structure formation. (Peaks compatible with Fig.~1 do not appear in non-baroque defect models, which now appear unlikely.) Though great strides have been made in the analysis of Boomerang and Maxima, there is intense needed effort worldwide now to develop new fast algorithms to deal with the looming megapixel datasets of LDBs and the satellites (\\eg Bond \\et 1999, Szapudi \\et 2000). ",
        "conclusions": ""
    },
    "0011/astro-ph0011470_arXiv.txt": {
        "abstract": "New results from the search for H$_2$ absorption in the damped Ly$\\alpha$ galaxy at redshift $z = 3.4$ toward QSO 0000--2620 ($z_{\\rm em} = 4.1$) are reported. The high-resolution ($\\lambda/\\Delta\\lambda = 48,000$) spectra of Q0000--2620 were obtained using the Ultraviolet - Visual Echelle Spectrograph (UVES) on the 8.2m {\\itshape ESO} Kueyen telescope. The ortho-H$_2$ column density is found to be $N(J=1) = (5.55 \\pm 1.35)\\times10^{13}$ cm$^{-2}$ ($2\\sigma$ C.L.). The combination of $N(J=1)$ with the limits available for other low rotational levels restricts the excitation temperature $T_{\\rm ex}$ in the range $(290 - 540)$~K. This gives the total H$_2$ column density of $N({\\rm H}_2) = (8.75 \\pm 1.25)\\times10^{13}$ cm$^{-2}$ and the corresponding fraction of hydrogen atoms bound in molecules of $f({\\rm H}_2) = (6.8 \\pm 2.0)\\times10^{-8}$. ",
        "introduction": "It has long been recognized that H$_2$ (and HD) molecules play a central role in the formation of gas condensations in the post-recombination era since they provide the cooling necessary for the collapse on all scales of the first objects. In the primordial gas at redshift $z \\la 50$, the fractional abundance of H$_2$ is calculated to be $f({\\rm H}_2) = 10^{-5} - 10^{-6}$ (e.g. \\cite{ref1.1}). The H$_2$ absorption lines from the Lyman and Werner bands may be observable at lower redshift $z \\sim 2 - 4$ in QSO absorption systems with high neutral hydrogen column densities, $N({\\rm HI}) = 10^{21} - 10^{22}$ cm$^{-2}$ (so-called Damped Lyman $\\alpha$ systems, DLA). It has been suggested that these systems are most closely related to the progenitors of normal galaxies~\\cite{ref1.2}. So far, absorption from H$_2$ has been detected in a few DLA systems which show, in general, a small amount of molecular gas~\\cite{ref1.2a}. In contrast to DLAs, observations in our Galaxy reveal the presence of H$_2$ in nearly all lines of sight in the disk and halo~\\cite{ref1.3}. It was also found that the ISM diffuse clouds with the neutral hydrogen column densities $N({\\rm HI}) \\ga 3\\times10^{20}$ cm$^{-2}$ show $f({\\rm H}_2) \\ga 10^{-5}$.  \\begin{figure} \\vspace{-1.5cm} \\centering \\includegraphics[width=1.0\\textwidth]{lev_fig1.eps} \\vspace{-1.0cm} \\caption[]{ Best fit to H$_2$ lines associated with the $z = 3.4$ DLA system toward Q0000--2620. ({\\bf a}, {\\bf b}, {\\bf c}, {\\bf d}) Simultaneous fit to the Lyman L(4-0)R(1), L(1-0)R(2) and Werner W(2-0)Q(1), W(2-0)R(0) lines with $N(J=0) = 8.4\\times10^{12}$ cm$^{-2}$ ($1\\sigma$ upper limit), $N(J=1) = (5.5 \\pm 0.7)\\times10^{13}$ cm$^{-2}$, $N(J=2) = 1.2\\times10^{13}$ cm$^{-2}$ ($1\\sigma$ upper limit), and $b \\simeq 10.07$ km~s$^{-1}$. The velocities shown are related to $z = 3.390127$. ({\\bf e}) Confidence regions for the ortho-H$_2$ column density calculated from the simultaneous fit of the L(4-0)R(1) + W(2-0)Q(1) lines and all low-ion lines located redward the Ly$\\alpha$ emission. ({\\bf f}, {\\bf g}) The corresponding deviations of the local continuum level. ({\\bf h}) The corresponding Doppler parameter variations. The grey areas restrict the $\\Delta C / C$ and $b$ values at $1\\sigma$ level in accordance with panel {\\bf e} } \\label{eps1} \\end{figure}  Absorption from H$_2$ in the $z = 3.4$ DLA system toward Q0000--2620 was not detected in the 1 \\AA\\, resolution spectrum obtained with the Multiple Mirror Telescope (MMT) and only an upper limit of $f({\\rm H}_2) < 3\\times10^{-6}$ has been reported previously~\\cite{ref1.4}. New observations with the UVES/VLT, having approximately 11 times higher resolution, revealed H$_2$ absorption in this DLA. Exactly at the expected position of the L(4-0)R(1) line where the MMT gave the limit $W_{\\rm rest} < 114$ m\\AA\\, (3$\\sigma$), an absorption line with the equivalent width $W_{\\rm rest} \\simeq 6$ m\\AA\\, was detected~\\cite{ref1.5}. The chance identification of this line in the Ly$\\alpha$ forest at $z \\sim 3$ has a probability $\\la 10^{-3}$ (cf.~\\cite{ref1.5a}).  Below we report on new identifications of H$_2$ lines in this DLA system and give improved values of molecular hydrogen column density and excitation temperature. ",
        "conclusions": ""
    },
    "0011/astro-ph0011536_arXiv.txt": {
        "abstract": "Interferometric techniques are at the forefront of modern astronomical instrumentation.  A new generation of instruments are either operating or nearing completion, including arrays of small telescopes as well as the ``big guns'' (VLTI and Keck).  A number of space interferometers for the detection of extra-solar planets are also being planned.  I will review the current state of play and describe the latest developments in the field. ",
        "introduction": "It is somewhat disappointing that a meeting on ``Galaxies and their Constituents at the Highest Angular Resolutions'' contains so few results from optical inter\\-ferometry\\footnote{In this review, the term ``optical'' is also meant to include infrared wavelengths.}.  To some extent, this reflects the interests of those on the Scientific Organising Committee, but it must also be taken as a sign that optical interferometry has yet to deliver on its promises.  Why is this?  One reason is surely the difficultly of the technology and its vulnerability to Hofstadter's Law (1980): \\begin{quotation} {\\em Hofstadter's Law} --- It always takes longer than you expect, even when you take into account Hofstadter's Law. \\end{quotation}  But progress is being made, and there is a large number of interferometry projects in various stages of development and operation.  The following is the current list of operating interferometers, listed in the order in which they obtained first fringes: \\begin{itemize}  \\itemsep=0pt \\item     GI2T/REGAIN (Grand Interf\\'erom\\`etre \\`a 2 T\\'elescopes). \\item     SUSI (Sydney University Stellar Interferometer) \\item     COAST (Cambridge Optical Aperture Synthesis Telescope) \\item     ISI (Infrared Spatial Interferometer) \\item     FLUOR (Fiber Linked Unit for Optical Recombination) \\item     IOTA (Infrared-Optical Telescope Array) \\item     NPOI (Navy Prototype Optical Interferometer) \\item     PTI (Palomar Testbed Interferometer) \\item     CHARA  (Center for High Angular Resolution Astronomy) \\end{itemize} The following are under construction (in alphabetical order, since I am not game to predict which will be first to get fringes!): \\begin{itemize}  \\itemsep=0pt \\item Keck Interferometer  \\item LBT (Large Binocular Telescope)  \\item MIRA-II (Mitaka IR Array)  \\item VLTI (Very Large Telescope Interferometer)  \\end{itemize} Finally, there are planned space missions, including SIM, TPF and DARWIN\\@.  More details of all these projects can be found via links from the Web-based ``Optical Long Baseline Interferometry News,'' currently maintained by Peter Lawson\\footnote{\\tt http://huey.jpl.nasa.gov/olbin/}. Here, I simply want to stress the number and range of the projects, and to point out that many of them are producing scientific results.  The proceedings of recent conferences (Unwin \\& Stachnik 1999; Lena \\& Quirrenbach 2000) testify to the vigorous activity in this field.  In Section~\\ref{sec.recent}, I highlight some areas in which recent technological developments have been made.  First, however, it is appropriate to make some general remarks about optical interferometry. ",
        "conclusions": ""
    },
    "0011/astro-ph0011299_arXiv.txt": {
        "abstract": "We reanalyse the UV/optical spectrum and optical broad-band data of the eclipsing binary HV 2274 in the LMC, and derive its distance following the method given by Guinan et al.~(1998a,b) of fitting theoretical spectra to the stars' UV/optical spectrum plus optical photometry.  We describe the method in detail, pointing out the various assumptions that have to be made; moreover, we discuss the systematic effects of using different sets of model atmospheres and different sets of optical photometric data. It turns out that different selections of the photometric data, the set of model atmospheres and the constraints on the value of the ratio of selective to total extinction in the $V$-band, result in a 25\\% range in distances (although some of these models have a large ${\\chi}^2$).  For our best choice of these quantities the derived value for the reddening to HV 2274 is $E(B-V)$ = 0.103 $\\pm$ 0.007, and the de-reddened distance modulus is DM = 18.46 $\\pm$ 0.06; the DM to the center of the LMC is found to be 18.42 $\\pm$ 0.07. This is significantly larger than the DM of 18.30 $\\pm$ 0.07 derived by Guinan et al. (1998a). ",
        "introduction": "The distance to the Large Magellanic Cloud (LMC) is a fundamental step in the cosmological distance ladder, since the extragalactic distance scale is usually determined with respect to the LMC distance. In fact, both the {\\sc hst} $H_0$ Key Project (Kennicutt et al. 1995, Freedman et al. 1999) and the Supernovae Calibration Team (Saha et al. 1999) fix the zero-point of the cosmological distance scale assuming a de-reddened LMC distance modulus (DM) of 18.50; in the case of the {\\sc hst} $H_0$ Key Project the adopted uncertainty of $\\pm$ 0.13 on the LMC distance modulus represents the largest contribution to their systematic error budget.  In recent years various methods have yielded LMC distance moduli showing a remarkable spread, ranging from DM = 18.07 (Udalski et al. 1998a) to DM = 18.70 (Feast \\& Catchpole 1997) -- see for example the review by Gibson (1999)  or Feast (2001). This uncertainty alone on the LMC distance causes an indetermination by $\\sim$20\\% on the value of the Hubble constant.     A method that, at present, seems to support the short distance scale, involves the analysis of the light-curve, radial velocity curve and UV/optical spectrum of the detached eclipsing binary HV 2274 in the LMC.  It is considered (see, e.g., Gibson 1999) to be one of the most promising techniques to derive a precise distance to the LMC, and it is based on a very elegant idea. From the analysis of the radial velocity and light-curve one obtains very accurate values for the masses and radii of the two binary components, as well as for the ratio of the effective temperatures. Fitting the UV/optical spectrum with model atmospheres one obtains the reddening, effective temperature and distance. This method, as already stressed, has been put forward as {\\sl the} way to obtain a very accurate DM to the LMC, in particular when more systems will be analysed.  Guinan et al. (1998a; hereafter G98a) found DM = 18.30 $\\pm$ 0.07 (the derived distance modulus to HV 2274 was 18.35 $\\pm$ 0.07. A geometric correction has been then applied to obtain the distance to the center of the LMC) after applying this technique.  More recently Nelson et al.~(2000) corrected the G98a value after re-determining the reddening toward HV 2274; they obtained DM = 18.40 $\\pm$ 0.07, a value only marginally in agreement with the long distance. However, they did not apply the method by G98a, deriving the distance correction only in an indirect way.  Because of the relevance of the method employed by G98a for deriving the LMC distance and in light of the recent claims by Nelson et al.~(2000) for a longer LMC distance, we want in the present paper to re-analyse the fitting procedure of the UV/optical spectrum, and carefully point out the uncertainties of this method. Unfortunately, in their 4-page {\\sl Letter}, G98a could not present all the intricate details that {\\sl are} involved in the application of this method, but which should be pointed out to the scientific community in order to judge the strengths and weaknesses of this technique (also see Feast 2001).  We will also study the sensitivity of the derived distance to both the set of model atmospheres and of optical photometric data employed in the fitting procedure.  In Sect.~2 we present an historical overview of observations and studies related to the HV 2274 distance. In Sect.~3 and 4, we discuss, respectively, the observational data and the method employed for the distance determination. Results are presented in Sect.~5, while a discussion follows in Sect.~6. ",
        "conclusions": "We will now discuss the results obtained in the previous section, with particular emphasis on the consistency of the derived parameters for the reference model.  \\subsection{The distance to HV 2274}  As previously explained, our reference model is model 5, which has the maximum wavelength coverage -- from the UV to the near infrared -- of the HV2274 spectrum, and makes use of the set of theoretical spectra which provides the lower $\\chi^2$ value.  In Fig.~1 and 2 we show, for model 5, the fit of the theoretical spectra to the FOS one, and the resulting UV and optical normalized extinction curve.  The final value for $\\left( \\frac{r_{\\rm A}}{d}\\right)^2$ is that of model 5. The final error estimate comes from the formal error in the model 5 result, added in quadrature to the squared differences of the parameter values of model 5 with those of models 7-16, and the additional error due to the error in  $r_{\\rm A}$. The result is $\\left( \\frac{r_{\\rm A}}{d}\\right)^2$ = (2.05 $\\pm$ 0.12) 10$^{-23}$. This corresponds to a linear distance of 49.10 $\\pm$ 1.44 kpc, or a true distance modulus of 18.46 $\\pm$ 0.06.     \\subsection{The distance to the center of the LMC}  To obtain the distance to the center of the LMC, the location of the binary with respect to the LMC center has to be taken into account. The geometry of the LMC can be described by an inclined disk, and one therefore has to consider the distance from the plane of that disk to the plane of the sky through the center of the LMC at the position of the binary, and the fact that the binary may be behind or in front of the plane of the disk.  G98a assume that HV 2274 is 1.1 kpc behind the center of the LMC, based on the parameters of Schmidt-Kaler \\& Gochermann (1992).  Various authors have described the geometry of the LMC with a thin disk and derived the position angle ($\\theta$) of the line-of-nodes and inclination angle ($i$). The relevant coordinate transformation to go from observed right ascension and declination to a rectangular coordinate system in the plane of the sky, and to a similar one rotated by $\\theta$ and $i$ are given in Weinberg \\& Nikolaev (2000). Table~6 gives estimates of the distance of the LMC plane to the center of the LMC for various accurate estimates of $\\theta$ and $i$, using the coordinate system by Weinberg \\& Nikolaev (2000). Because of the different orientations and definition of positive inclination the values in the table may differ by 90\\degr\\ with respect to the values quoted in the original reference.  The mean of these 5 determinations is 0.88 kpc; the average of the highest and lowest value is 0.89 kpc and the median is 0.82 kpc. The adopted difference in distance is 0.9 $\\pm$ 0.3 kpc\\footnote{As an aside we did the same for SN 1987A. The shortest distance (SN 1987A - LMC Center) is $-0.64$ kpc [for the parameters of Schmidt-Kaler \\& Gochermann (1992)], and the largest is $-0.15$ kpc [for the parameters of Groenewegen (2000)]. Based on all 5 determinations, the LMC plane at the location of SN1987A is 0.4 $\\pm$ 0.2 kpc in front of the LMC center.}.  The error does not come from the internal errors of each of the determinations, but from the spread among the values itself, and the difference between the lowest and highest value has been assumed to correspond to 3$\\sigma$.  It is unknown if HV 2274 is in front or behind the LMC plane. The vertical scale height of the LMC disk is small however, between 100 and 300 pc (see the discussion in Groenewegen, 2000). The largest value is taken here, and added to the error mentioned above, to give the final result that HV 2274 is located 0.9 $\\pm$ 0.5 kpc behind the LMC center.  Considering a DM to HV 2274 of 18.46 $\\pm$ 0.06, the DM to the LMC center 18.42 $\\pm$ 0.07.    \\begin{table} \\caption[]{The distance (HV 2274 $-$ LMC center) }  \\begin{tabular}{cccccccc}  \\hline $\\theta$ & $i$     & Reference & $\\Delta$ \\\\ (\\degr) & (\\degr) &           & (kpc)    \\\\ \\hline 258     & 38  & Schmidt-Kaler \\& Gochermann (1992) & 1.30 \\\\ 258     & 33  & Feitzinger et al. (1977)           & 1.08 \\\\ 261     & 25  & Weinberg \\& Nikolaev (2000)        & 0.82 \\\\ 296     & 18  & Groenewegen (2000)                 & 0.75 \\\\ 232     & 29  & Martin et al. (1979)               & 0.47 \\\\ \\hline \\end{tabular} \\end{table}   \\subsection{The extinction curve towards HV 2274}  From model 5 we get $x_0 = 4.67 \\pm 0.04$, $\\gamma = 1.27 \\pm 016$, $c_3 = 2.03 \\pm 0.58$, $c_4 = 0.66 \\pm 0.10$ and $R = 3.06 \\pm 0.14$ (where the errors are the internal errors due to the fitting procedure only); from Eqs.~(6) and (7) one obtains $c_1 = -0.13 \\pm 0.21$ and $c_2 = 0.72 \\pm 0.07$.  These values are consistent with the corresponding quantitites recently determined by Misselt et al.~(1999) for their 'LMC-Average Sample' of stars.  Figure 2 shows the resulting normalised UV and optical extinction curve.    \\subsection{The reddening towards HV 2274}  The value of $E(B-V)$ obtained from model 5 is $E(B-V)$ = 0.103 $\\pm$ 0.007.  It is important to notice that this value is different from the one derived by G98b ($E(B-V)$ = 0.120 $\\pm$ 0.009) and from the determination by UPW98 ($E(B-V)$ = 0.149 $\\pm$ 0.015) who used colour-colour $(U-B)-(B-V)$ relationships.  Since we are using the $(B-V)$ (and $(V-I)$) photometric data by UPW98 as constraint for the spectrum fit we investigated the possibility that the reddening derived from the spectrum fit is inconsistent with HV2274 broad band photometry.  We followed the same procedure by UPW98 that is, to consider a local standard $(U-B)-(B-V)$ sequence of stars with the same spectral type as HV 2274, and derive the reddening from the displacement of the position of HV 2274 with respect to the standard sequence, for an assumed $E(U-B)/E(B-V)$ ratio. We used the same standard colour-colour sequence employed by UPW98 (for B-stars of luminosity class III), and a ratio $E(U-B)/E(B-V)$ derived from our adopted extinction law using the value of $R$ derived from the fit (which is by the way very close to the standard value of 3.1).  By computing appropriate stellar models (using the same input physics and colour transformations as in Salaris \\& Weiss 1998 and a scaled-solar metal distribution) for solar metallicity and the metallicity derived from model 5, with masses around 12$M_{\\odot}$, we verified that the location of B stars on the $(U-B)-(B-V)$ plane does not vary in this metallicity range.  As for the $(B-V)$ colour of HV 2274 we used the value by UPW98, while for $(U-B)$ we derived the value from the FOS spectrum (see Section 4.2), since it covers the wavelength region spanned by this colour index.  The $(U-B)$ colour derived from the spectrum is different from the value observed by UPW98.  The FOS spectrum provides $(U-B)$ = $-$0.836 $\\pm$ 0.006, while UPW98 measured $(U-B)$ = $-$0.905 $\\pm$ 0.04. Nelson et al.~(2000) measured $(U-B)$ = $-0.793$ $\\pm$ 0.031 which is closer (but still inconsistent at the 1$\\sigma$ level) to the value derived from the FOS spectrum. From the colour-colour diagram we get $E(B-V)$ = 0.115 $\\pm$ 0.015 which is consistent, within 1$\\sigma$, with the value derived from model 5.   \\subsection{The metallicity of HV 2274}  The value of [m/H] from the reference model, taking the formal error from the fit plus the external errors from models 7-16 is $-0.38 \\pm 0.12$. This is relative to the adopted iron abundance of 7.67 in the Kurucz models. This implies that the metallicity relative to the currently favoured solar iron abundance of 7.51 (Grevesse \\& Noels 1993) is [Fe/H] = $-0.22 \\pm 0.12$. Model 6 computed using Butler atmospheres has [m/H] = $-0.13$ relative to a solar abundance of 7.46, [Fe/H] = $-0.18$ relative to the solar abundance of Grevesse \\& Noels. These values are in good agreement with each other, but significantly more metal rich than the value of $-0.42 \\pm 0.07$ to $-0.45 \\pm 0.06$ found by G98a,b.  Observationally there are few direct iron abundance determinations for hot and young stars in the LMC. Haser et al. (1998) derived a metallicity of $-0.3$ and $-0.1$ dex for an O3{\\sc iii} and O4{\\sc i} star in the LMC. Korn et al. (2000) determined the iron abundance in 5 non-supergiant B-stars, the average value being $-0.42 \\pm 0.15$. Both the values found by G98a,b and ours are consistent with these determinations.   \\subsection{The \\M\\ data}  As mentioned in Sect.~3 HV 2274 has been detected in the \\M\\ $JHK_{\\rm s}$ infrared survey. The observation date is JD = 2451111.689. From the ephemeris in Watson et al. (1992) a primary eclipse is predicted at JD = 2451111.697 $\\pm$ 0.006, taking into account the error in the time determination of the reference primary eclipse and the period. The shift of 0.008 $\\pm$ 0.006 days, or 0.0014 in phase is negligible, also in light of the fact that phase shifts of up to 0.05 occur due to apsidal rotation (Watson et al. 1992).  In the optical the magnitude difference between out-of-eclipse and primary eclipse are 0.72, 0.71 and 0.70 mag in $BVI$, respectively (Watson et al. 1992). Naively, one might therefore expect a magnitude difference of about 0.69 in $J$.  The difference between the observed \\M\\ $J$ and the model predictions is (15.152 $\\pm$ 0.057) $-$ (14.503 $\\pm$ 0.049) = 0.65 $\\pm$ 0.08, where the error in the model comes from summing up all the error terms from models 6-17. This is in good agreement with the value extrapolated from the magnitude difference between out-of-eclipse and primary eclipse measured in the optical. However, as emphasised later, out-of-eclipse NIR photometry would be important in further constraining the parameters of the model.   \\begin{figure} \\centerline{\\psfig{figure=h2475f2.ps,width=8.5cm}} \\caption[]{UV and optical normalised extinction curve for the parameters of model 5.} \\end{figure}  \\subsection{Final remarks}  In this paper we have described and analyzed in detail the method used by G98a,b to derive the distance to the eclipsing binary HV 2274 in the LMC. We used various sets of theoretical spectra and broad band photometric data in the fitting procedure outlined in Section 4, and we found that Kurucz ATLAS 9 spectra best reproduce the observed spectrum of HV 2274.  The selection of the wavelength range to be covered by the spectrum fit and the constraint on $R$ play also a role in determining the outcome of the fitting procedure. We are now going to comment briefly about this point.  In the reference model we fitted HV 2274 FOS data plus $(B-V)$ and $(V-I)$ colours (from UPW98) using ATLAS 9 spectra, constraining the parameters $c_1$ and $c_2$ of the UV extinction law according to Equations (6) and (7) and keeping $R$ as a free parameter.  The use of $(V-I)$ is dictated by our desire to use all available information about the spectral energy distribution of HV 2274.  The homogeneous $UBVI$ photometry by UPW98 makes it possible to cover all the spectral range from UV to near-infrared.  Neglecting $(V-I)$, that means, fitting a smaller wavelength range, induces a decrease by 0.10 mag in the distance modulus derived from the fitting procedure.  If we consider the ATLAS 9 models as an accurate reproduction of the 'real' spectra of B stars, this difference can be ascribed to the fact that there is information about the distance contained in the $(V-I)$ colour, and therefore it must be included in the fitting procedure.  Conversely, this result could also imply that ATLAS 9 spectra are inconsistent with observations at the longer wavelengths or that broad-band photometry and FOS spectrum are not homogeneously calibrated. In this respect, one should keep in mind that the $(U-B)$ value derived from the FOS spectrum is inconsistent with the $(U-B)$ determined by UPW98 and Nelson et al.~(2000) and also that the $(U-B)$ values derived from the best-fit ATLAS 9 spectrum are systematically different by $\\sim$0.03 mag with respect to the FOS colour (see Table~5).  In case of not considering $(V-I)$ for the spectrum fitting, keeping $R$ fixed increases sensibly (0.09 mag) the derived distance modulus with respect to the case of having $R$ determined by the fitting procedure, as shown by comparing the outcome of models 21 and 17. However, it is not clear why one should fix the value of $R$ a priori, since it is known that it is subjected to variations within the LMC and within our galaxy.  Finally, the use of the $(B-V)$ by Nelson et al.~(2000) together with the FOS spectrum produces a distance modulus even larger (0.16 mag larger than the case of using UPW98 data).  The quality of the fit is only marginally lower, and the reddening compares well with the value ($\\sim$ 0.08) one would derive from the colour-colour diagram considering the Nelson et al.~(2000) $(B-V)$ and the FOS $(U-B)$.  A last comment is that it is very unfortunate that the \\M\\ infrared data happens to be taken during eclipse. The difference in predicted $JHK$ between the different models is up to 0.15 mag (see Table~5). Therefore, accurate NIR photometry or even NIR spectroscopy at the 1\\% level is expected to give valuable additional constraints. \\\\  The conclusion is that the use of eclipsing binaries as distance indicators is powerful, but that a photometrically well calibrated data set covering a large wavelength region is essential. Furthermore, the discrepancies among the different theoretical model atmospheres is worrying and need further investigation."
    },
    "0011/astro-ph0011250_arXiv.txt": {
        "abstract": "We explore the relative importance of the stellar mass density as compared to the inner dark halo, for the observed gas kinematics thoughout the disks of spiral galaxies. We perform hydrodynamical simulations of the gas flow in a sequence of potentials with varying the stellar contribution to the total potential. The stellar portion of the potential was derived empirically from K-band photometry. The output of the simulations -- namely the gas density and the gas velocity field -- are then compared to the observed spiral arm morphology and the H$\\alpha$ gas kinematics. We solve for the best matching spiral pattern speed and draw conclusions on how massive the stellar disk can be at most. For the case of the galaxy NGC 4254 (Messier 99) we demonstrate that the prominent spiral arms of the stellar component would overpredict the non-circular gas motions unless an axisymmetric dark halo component adds significantly in the radial range ${\\rm R_{exp} < R < 3\\, R_{exp}}$. ",
        "introduction": "In almost all galaxy formation scenarios, non-baryonic dark matter plays an important role. Today's numerical simulations of cosmological structure evolution quite successfully reproduce the observed galaxy distribution in the universe \\cite{kau}. While galaxies form and evolve inside dark halos their physical appearance depends strongly on the local star formation and merging history. At the same time the halos evolve and merge as well. According to the simulations, we expect that the dark matter is important in the inner parts of galaxies \\cite{NFW1},\\cite{NFW2} and that it thus has a considerable influence on the kinematics. These predictions are in contrast to some studies which indicate that galactic stellar disks - at least of barred spiral galaxies -  alone dominate the kinematics of the inner regions \\cite{deb}. Apparently this is also the case in our own Milky Way \\cite{ger}. \\\\ Determining individual mass fractions of the luminous and dark matter is not a straightforward task. The rotation curve of a disk galaxy is only sensitive to the total amount of gravitating matter, but does not allow the distinction between the two mass density profiles. For a detailed analysis it is necessary to adopt more refined methods to separate out the different profiles. Previous investigations used for example knowledge of the kinematics of rotating bars \\cite{wei} or the geometry of gravitational lens systems \\cite{mal}. Here we would like to exploit the fact, that the stellar mass in disk galaxies is often organized in spiral arms, i.e.~in coherent non-axisymmetric structures. In most proposed scenarios, the dark matter, however, is collisionless and dominated by random motions. Therefore it is not susceptible to spiral structures. If the stellar mass dominates, the spiral arms, as traced by the near infrared (NIR) light, should induce considerable non-circular motions in the gas, that manifest themselves as velocity ''wiggles'' in observed gas kinematics. Using hydrodynamical gas simulations we are able to predict these velocity wiggles and compare them to the observations. Hence the contribution of the perturbative forces with respect to the total forces can be determined quantitatively and can be used to constrain the stellar disk to halo mass ratio. ",
        "conclusions": "Although there is quite some scatter in the observed gas kinematics, we find that the velocity jumps, which are apparent in the simulations for the maximum disk case are too large to be in agreement with the measurements. The inner part of the simulated rotation curve ($<$ 0\\farcm3) is dominated by the dynamics of the small bar, which is present at the center of the galaxy. Its pattern speed might be different from the one of the spiral's and thus relate to a mismatch in the inner part of the rotation curve. We conclude that an axisymmetric dark halo is needed to explain the kinematics of the stellar disk. The influence of the stellar disk is submaximal in the sense that we don't find strong enough velocity wiggles in the observed kinematics as would be expected if the stellar disk was the major gravitating source inside the inner few disk scale lengths. How this conclusion might apply to other spiral galaxies will be the upcoming issue of this project. We plan to extend our analysis at first to the 3 other galaxies where we have already now complete data sets. Finally we intend to draw our final conclusions on a basis of a sample consisting of 8 - 10 members. This should be sufficient to determine reliable results about the luminous and dark mass distributions in spiral galaxies."
    },
    "0011/astro-ph0011066_arXiv.txt": {
        "abstract": "Deep submillimetre surveys have successfully detected distant, star-forming galaxies, enshrouded in vast quantities of dust and which emit most of their energy at far infrared wavelengths. These luminous galaxies are an important constituent of the Universal star-formation history, and any complete model of galaxy evolution must account for their existence. Although these sources have been tentatively identified with very faint and sometimes very red optical counterparts, their poorly constrained redshift distribution has made their interpretation unclear. In particular, it was not understood if these galaxies had been missed in previous surveys, or if they constituted a truly new class of objects, undetectable at other wavelengths.  By utilizing a radio selection technique, we have isolated a % sample of 20 sub-mm objects representative of the 850\\mum\\ population brighter than 5\\,mJy with $z$\\cle 3. We show that these galaxies are so heavily dust obscured that they remain essentially 'invisible' to ultraviolet selection. Furthermore, relying on the radio-submillimeter flux density ratio, we estimate their redshift distribution, finding a median of two. These results are inconsistent with the existence of a very high redshift ($z>4$) population of primeval galaxies (L$_{bol} > 10^{12}$\\,h$^{-2}$\\Lsun) contributing substantially to the sub-mm counts. While not a substitute for the thorough followup of blank field sub-mm surveys, our results do shed light on a substantial portion of the luminous sub-mm population with $z$\\cle 3. ",
        "introduction": "The extragalactic far-infrared background light is believed to be composed of the integrated thermal starlight and non-thermal AGN radiation, reradiated by dust within star-forming galaxies over the entire history of galaxy formation. The energy density of this infrared background is approximately the same as found in the optical suggesting that at least half of the Universal star-formation history remains hidden from optical view (Puget et al.~1996). This diffuse background was first resolved into discrete sources by the Sub-millimetre Common User Bolometer Array (SCUBA -- Holland et al.~1999) on the James Clerk Maxwell Telescope (JCMT) by a number of groups (Smail et al.~1997, Hughes et al.~1998, Barger et al.~1998, Eales et al.~1999).  Although a large number of deep SCUBA surveys has led to a better estimate of the 850 micron galaxy surface density, our understanding of the nature of the sub-mm population remains limited. The principal obstacle is obtaining reliable counterparts of these sub-mm sources at other wavelengths, a problem exacerbated by both the coarseness of the JCMT resolution (15 arcsec at 850 microns) and the inherent faintness of suspected optical counterparts (Smail et al.~1999). It is still unclear if the sub-mm selected sources are related to known populations, such as high redshift quasars (Hughes et al.~1997, McMahon et al.~1999) or Lyman-break galaxies (Chapman et al.~2000a), or constitute a truly new class of objects. With the exception of several isolated objects, few reliable identifications have been made (e.g. Ivison et al.~1998, Frayer et al.~1998). Thus the redshift distribution has remained largely unconstrained over a vast range, with the possibility that many sources lay at extreme distances ($z > 5$).  One technique, which has shown some promise in identifying sub-mm sources, is radio continuum followup. Because galaxies and the inter-galactic medium are transparent at centimeter wavelengths, radio emission is unhindered by intervening gas and dust. Ubiquitous in local star-forming galaxies, radio emission also correlates very strongly with the far-infrared emission in star-forming galaxies (Helou et al.~1986). Moreover, the high resolution provided by radio interferometers can provide a surrogate for the poor sub-mm angular resolution and astrometric uncertainties.  Given the difficulties of obtaining secure optical identifications and spectroscopic redshifts, the radio observations provide another clue to the nature of sub-mm sources. Via the empirically observed far-infrared to radio correlation in local star-forming galaxies, one can use the observed ratio of sub-mm to radio continuum flux density to estimate a redshift.  As the k-corrections (corrections based on the redshifted spectral energy distribution - SED) of the radio and sub-mm flux densities are opposite in slope, the ratio of radio to sub-mm flux density is quite sensitive to redshift (Carilli \\& Yun 1999).  Barger, Cowie \\& Richards (2000 -- hereafter BCR) first attempted to use a radio selected sample to target a number of optically faint microJansky radio sources with near-infrared magnitude, $K > 20.5$. Using SCUBA to a 3$\\sigma$ RMS limiting flux density of 6\\,mJy at 850 microns, they detected 5 out of 15 radio sources, while in the process demonstrating that none of the optically brighter radio sources ($K < 20.5$) were detected in the sub-mm. The surface density of these few bright radio selected sub-mm sources closely matched that from blank field surveys, indicating a close correspondence between the optically faint radio population and bright sub-mm sources. Other pointed SCUBA studies of known high-$z$ populations such as $z\\sim3$ Lyman-break galaxies (LBGs -- Chapman et al.~2000a) and radio loud AGN (Archibald et al.~2000) have revealed few SCUBA detections, and nowhere near the surface density of blank field sub-mm sources.  We have refined the selection criterion to those microJansky radio sources with an optical magnitude, $I > 25$, based on the clear bi-modal break in the optical properties of microJy radio sources (Richards et al.~1999). We have applied this technique to a sample in the region surveyed by Richards (2000) centered on the Hubble Deep Field. We have selected a total of 47 radio sources in our study, 20 previously observed, which meet our criterion. Our followup SCUBA {\\it photometry} observations of 27 radio selected objects demonstrate this to be a highly successful technique. We now detect $\\sim$50\\% of the new 27 object sample above 4.5\\,mJy at 850 microns, with an overall success rate of 20 out of 47 objects observed. Our increased detection success over BCR is likely to be a result of our slightly deeper survey coupled with the stricter selection criterion. We are thus able to uncover bright sub-mm sources using SCUBA at the rate of one source per hour on the JCMT, greater than an order of magnitude more rapid than mapping a random patch of sky. Our new survey represents a sub-mm mapping of a $\\sim$100 arcmin$^2$ effective region in less than 16 hours, sensitive to sources S$_{850}>5$\\,mJy and $z$\\cle 3. ",
        "conclusions": "The crucial data available to us from our technique are the optical properties and redshift estimates for the sub-mm sources, which we present in Table~1.  Our results assume a $\\Lambda=0.0$, $\\Omega=1.0$, H$_0$=65\\,km/s/Mpc cosmology. Redshift estimates can be obtained from the sub-mm/radio flux ratios (Carilli \\& Yun 2000). All of our sub-mm sources fall roughly in the redshift range $z=1-3$ with a median redshift for the sample of $z=1.9$, consistent with previous results from BCR. The sensitivity of our radio survey to star-forming galaxies with radio luminosities fainter than $10^{24}$\\,W/Hz diminishes quickly past $z\\sim3$, and hence biases our sub-mm survey. An independent check on the sub-mm/radio redshift estimates can be obtained through the 450\\mum/850\\mum\\ ratio (e.g. Hughes et al.~1998). Subject to unknown dust temperature, we obtain an estimate of T$_{\\rm d}$/(1+$z$), which we list in Table~1 for T$_{\\rm d}$=45\\,K for consistency with the ultra-luminous infrared galaxy, Arp\\,220. Raising or lowering the adopted dust temperature has the effect of a corresponding systematic raising and lowering of both our redshift estimates (Blain 1999).  In order to calculate the density of sub-mm sources on the sky as found in our radio pre-selection survey, we need to determine the effective area covered by our study. This is given simply by the overlap regions between the optical and radio images. However, a further complication arises from the non-uniform sensitivity of the radio image which serves to decrease the visibility area. This issue is discussed in Richards (2000) and we use the same method for determining the source count.  Our pre-selected sample has already saturated the bright ($>5$\\,mJy) sub-mm counts (Fig.~1), % and there are not likely to be many additional bright, high redshift sub-mm sources in our survey region. BCR found that 2 additional sub-mm sources without radio counterparts were detected in their survey area, indicating 75\\% of the bright sub-mm sources are typically recovered through such radio pre-selection.  Since we performed photometry on the sources, we have no means of estimating this extra population. By accounting for such an additional 25\\% high redshift population our sample is in agreement with analyses of the redshifts for lensed sub-mm sources (Barger et al.~1999, Smail et al.~2000, Blain et al.~1999b). The percentage of sub-mm sources missed by our pre-selection technique will depend field to field on the high-$z$ clustering of sub-mm luminous sources. A lensed sub-mm survey (Smail et al.~2000) detects a similarly large fraction of their bright sub-mm sources in the radio. Deeper blank field sub-mm surveys (e.g. Eales et al.~2000 -- S$_{850\\mu m}$\\cge 3\\,mJy) detect $\\sim$1/3 of their sources in the radio. This is roughly in agreement with our results which use deeper radio maps and brighter sub-mm limits. While not a substitute for the thorough (and difficult) followup of blank field sub-mm surveys, our results do shed light on a substantial portion of the luminous sub-mm population with $z$\\cle 3.  At the faint end of the counts, our pre-selected sources appear to fall short of full source counts recovery, and the interesting question becomes what is the nature of the 850-micron sources at flux densities $<$5\\,mJy? Although we are only sensitive to sources with S$_{850}$\\cge 5\\,mJy, we can deduce important properties about the fainter sub-mm population. Averaging our sub-mm undetected sample (inverse variance weighted) reveals a mean flux of S$_{850}=0.8\\pm0.3$\\,mJy, suggesting that many of these $\\sim$50\\% of our $I>25$ and S$_{1.4}>40$\\,$\\mu$Jy radio sample are still fairly luminous sub-mm sources. They likely form a continuous distribution with the S$_{850}$\\cge 5\\,mJy sample, lie at similar or lower redshifts (Table~1), and comprise roughly 10\\% of the blank field SCUBA source counts from 1--5\\,mJy.  This leaves a large portion of sub-mm sources fainter than 5\\,mJy that are not subsumed in our present $I>25$ radio sample.  The first possibility is that these fainter sub-mm sources are largely contained in our present radio surveys, but are actually optically bright ($I<25$).  Bright LBGs are known to emit at the 1--2\\,mJy level (Chapman et al.~2000a, Peacock et al.~2000), so they must contribute some fraction of the missing sources.  This is consistent with the high SFR deduced by Steidel et al.~1999, who apply a large dust correction to their results.  Our Arp\\,220 SED model suggests that these sources would be detected in our present radio survey almost out to $z=3$.  However, the HDF-SCUBA results (Hughes et al.~1998, Peacock et al.~2000) also show directly that a significant fraction of the 2\\,mJy sources are not associated with bright LBGs.  Although deeper radio observations with optically faint counterparts may quickly recover this population, it also remains a possibility that these sources represent high redshift ($z>4$) protogalaxies with L$_{\\rm bol} < 10^{12}$\\Lsun, which would remain undetected in the radio to significantly deeper flux limits. So by pushing to fainter radio limits, it is rather unclear what types of objects might be pre-selected.  An additional concern with the sub-mm source population is that massive AGN may be heating the dust rather than star formation. We can be reasonably certain that our radio selected sub-mm sources are not driven primarily by AGN for two reasons. Firstly, the sources are spatially resolved with a median of about 2\\arcsec\\ in the radio using the Merlin interferometer at a resolution of 0.2\\arcsec\\ (Richards 2000, Muxlow et al.~2000), corresponding to $\\sim$1\\,kpc at $z=2$ for our adopted cosmology. If the radio and associated sub-mm emission were emanating from such a compact active nucleus (AGN), it would appear unresolved even at this fine resolution.  Secondly, the sub-mm sources have recently been shown to have little or no overlap with X--ray sources as observed with the Chandra satellite (e.g. Fabian et al.~2000, Hornschemeier et al.~2000, Barger et al.~2001). As scenarios which would obscure even the X--ray emission are improbable, the implication is that most bright sub-mm sources are in fact driven by star formation.  Assuming then that our radio-selected sources are driven primarily by star formation, it is appropriate to use our data to estimate a contribution to the comoving star formation rate density (SFRD).  In Fig.~2, we integrate over our measured counts in three redshift bins, and divide by the effective volume of our detected sources, to represent the sub-mm SFRD as a function of redshift. Since we have assumed the far-IR/radio relation in the redshift estimates, our SFRD estimates can be derived from either wavelength in a manner similar to BCR. We plot our new points as solid squares (sub-mm detected objects) and a solid hexagon (sub-mm undetected objects).  The very high redshift open square ($z>3$) represents the Hubble Flanking Fields sub-mm sources undetected in the radio from BCR and Borys et al.~2001. We then compare optically selected sources (open triangles) at redshifts $1<z<4$ (Connolly et al.~1997, Steidel et al.~1999). We apply a dust correction to the optical points for $z>1$, using the Steidel et al.~(1999) prescription (factor $\\sim$5 for $z>2$). This is roughly in accord with the expected correction from the Chapman et al.~(2000b) sub-mm measurement of 33 LBGs with large expected star formation rates.  Lower redshift radio selected sources with bright optical counterparts, as analysed by Haarsma et al.~(2000) are also plotted (red open circles) requiring no correction factor for dust obscuration.  Sub-mm sources fainter than 5\\,mJy likely begin to merge with optically selected samples (e.g. Adelberger \\& Steidel 2000).  However, our $I>25$ radio selected population is truly an orthogonal population to those discovered in optical surveys, even for S$_{850}$\\cle 5\\,mJy. We can then confidently sum the optically selected (but uncorrected for dust) ($I<25$) and sub-mm ($I>25$) points from $z=$1--4 to arrive at a conservative lower limit to the total SFRD of all presently known objects (plotted as large stars with arrows). This lower limit can be compared with the dust corrected optical points, which is still a rather uncertain procedure.  We have therefore recovered the majority of the bright ($>5$\\,mJy) sub-mm sources with a statistically significant sample, selected based on the microJy radio emission with extremely faint optical counterparts. We can state with some assurance that the bulk of bright sub-mm selected sources represent a highly dust obscured star forming population which would be very difficult to identify in optically based surveys. Our redshift analysis has also demonstrated that this is not because the sub-mm sources are at very high redshifts ($z>>3$). These sources may represent the epoch in which the most massive spheroid galaxies were being formed through merging fragments in cluster environments. Indeed, the recent identification of prodigious sub-mm emission with highly over-dense cluster cores at $z\\sim3$ and $z\\sim3.8$ (Chapman et al.~2000c and Ivison et al.~2000 respectively) suggests that the most luminous members of our sub-mm source sample may highlight similar such regions.  Pushing our study to fainter sub-mm and radio flux limits will facilitate our understanding of the transition and overlap between these ultra-luminous high$-z$ star formers (which may evolve into the most massive spheroids in the present epoch) and the less massive galaxies selected in the optical through the Lyman-break technique (Steidel et al.~1999)."
    },
    "0011/astro-ph0011316_arXiv.txt": {
        "abstract": "\\noindent We present analytic solutions of Maxwell equations in the internal and external background spacetime of a slowly rotating magnetized neutron star. The star is considered isolated and in vacuum, with a dipolar magnetic field not aligned with the axis of rotation. With respect to a flat spacetime solution, general relativity introduces corrections related both to the monopolar and the dipolar parts of the gravitational field. In particular, we show that in the case of infinite electrical conductivity general relativistic corrections due to the dragging of reference frames are present, but only in the expression for the electric field.  In the case of finite electrical conductivity, however, corrections due both to the spacetime curvature and to the dragging of reference frames are shown to be present in the induction equation. These corrections could be relevant for the evolution of the magnetic fields of pulsars and magnetars.  The solutions found, while obtained through some simplifying assumption, reflect a rather general physical configuration and could therefore be used in a variety of astrophysical situations. ",
        "introduction": "The investigation of the influence of strongly curved spacetimes on the properties of electromagnetic fields has an interest of its own which is further increased when these effects could be related to a rich observable phenomenology. This coupling between general relativistic effects and electromagnetic fields is expected to be particularly important in the vicinity of neutron stars which are among the most relativistic astrophysical objects and are characterized by very intense magnetic fields (Lamb 1991, Glendenning 1996). A number of different observations indicate that in young neutron stars the surface magnetic field strengths are of the order of \\hbox{$10^{11}-10^{13}$ G.}  In some exceptional cases, as those of magnetars, magnetic field strengths \\hbox{$\\ge 5 \\times 10^{14}$ G} are considered responsible for the phenomenology observed in soft gamma-ray repeaters (Duncan \\& Thompson 1992, Thompson \\& Duncan 1995). Older neutron stars, observed as recycled pulsars and low mass X-ray binaries, show instead surface magnetic fields that are much weaker $\\le 10^{10}$ G suggesting that these are subject to a decay, even if it is still difficult to establish whether the decay is due to accretion (Geppert \\& Urpin, 1994; Konar \\& Bhattacharya, 1997) or to other processes.  In the case of isolated neutron stars, the possibility of magnetic field decay as a result of accretion does not arise, but there are still a number of different ways in which the energy stored in the magnetic can be lost. This can happen either through the emission of electromagnetic (dipole) radiation, through Ohmic decay, through ambipolar diffusion, or through more complicated effects such as ``Hall cascades'' (see Goldreich and Reisenneger 1992 for a review). The investigation of these scenarios requires combined efforts. On one hand, there is the search for a more precise description of the microphysics of the processes involved, some of which are still not well quantified. On the other hand, attention is paid to a more realistic description of the gravitational effects on the properties of the electromagnetic fields in highly curved spacetimes and this is also the motivation of this work.  The investigation of the general relativistic corrections to the solution of Maxwell equations in the spacetime of a relativistic star has a long history. The initial works of Ginzburg \\& Ozernoy (1964), Anderson \\& Cohen (1970) and of Petterson (1974) on the stationary electromagnetic fields in a Schwarzschild spacetime have revealed that the spacetime curvature produces magnetic fields which are generally stronger than their Newtonian counterparts (see also Wasserman \\& Shapiro 1983 for a subsequent derivation). Sengupta (1995) has reconsidered this problem and also looked for a general relativistic expression for the electric field in the Schwarzschild background of a neutron star. As we will discuss in Section \\ref{srst_es} the method used in his derivation is not entirely correct and the results obtained for the electric field are not solutions of Maxwell equations. More recently, Sengupta has also considered the problem of the Ohmic decay rate in a Schwarzschild spacetime (Sengupta, 1997). His approach is strictly valid only for the region of spacetime external to the star as it does not provide a correct general relativistic description of the electromagnetic fields internal to the star. Within these approximations, however, Sengupta (1997) has pointed out that the effects of intense gravitational field seem to decrease the overall decay rate by a couple of orders of magnitude. The same problem has also been considered in more detail by Geppert, Page and Zannias (2000). Their analysis was aimed at a mathematically consistent solution of Maxwell equations also in the spacetime region internal to the star and makes therefore use of a generic metric for a non-rotating relativistic star. Their results, while confirming a decrease in the typical decay time for the magnetic field, also show that the decay time is smaller but comparable with the one found in flat spacetime.  The general relativistic effects induced by the rotation of the star were first investigated by Muslimov \\& Tsygan (1992) in the slow rotation approximation. A similar approach was also used by Muslimov and Harding (1997) for the electromagnetic fields external to a rotating magnetized star. Their analysis refers to a charge filled magnetosphere and represents the relativistic extension of the Goldreich-Julian model. Using a different derivation, Prasanna and Gupta (1997) have also investigated the properties of the electromagnetic fields in the magnetosphere of a relativistic rotating neutron star, with special attention being paid to the dynamics of charged test particles.  We here extend and unify all of the above investigations by considering the solution of Maxwell equations in the internal and external background spacetime of a slowly rotating magnetized relativistic star. The star is considered isolated and in vacuum, with a dipolar magnetic field which is not assumed aligned with the axis of rotation. The purpose of this paper is threefold.  Firstly, we want to extend previous results to the most general case of a misaligned rotator, providing for this case also the form of the electric field. Secondly, we want to discuss the possible role played by frame dragging effects in the Ohmic decay for an isolated neutron star and estimate its importance. Thirdly, we wish to clarify a few important aspects of the solution of Maxwell equations in the gravitational field of a relativistic star that, when overlooked, have led to incorrect solutions (Sengupta 1995, Prasanna and Gupta 1997). Finally, by providing a rather general solution to the problem (although truncated at the lowest order in the expansion of the angular dependence) we offer a compact reference from which all of the previous results can be easily found in the appropriate limits and which could have practical astrophysical applications.  The paper is organized as follows: in Section \\ref{meq} we write the general relativistic Maxwell equations in the metric of a slowly rotating star and the form they assume when the electromagnetic fields are those measured in the orthonormal frame of zero angular momentum observers. In Section \\ref{ss} we find the stationary solutions (i.e. solutions in which the infinite conductivity of the medium prevents a variation in time of the star's magnetic moment) to Maxwell equations outside and inside the misaligned rotating star. For this we consider first the problem in Newtonian gravity and we then extend the results to general relativity within the slow rotation approximation. Section \\ref{nss} is devoted to the equivalent problem, but in the case in which the magnetic field is not supposed stationary. There, we derive the basic induction equations for the evolution of the inner stellar magnetic field of a misaligned rotating star. Section \\ref{conclusion} contains our conclusions and the prospects of future developments.  A number of appendices provide further details about some of the calculations carried out in the main part of the paper. In particular, Appendix A summarizes the components of the electromagnetic tensor in a coordinate basis and in a locally orthonormal tetrad, while Appendix B shows the derivation of the radial eigenfunctions for the electromagnetic fields in terms of Legendre's equation. Appendix C shows the explicit expressions for the surface charges and currents and, finally, Appendix D contains an alternative and equivalent derivation of the equations for the time evolution of magnetic field in terms a vector potential. Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = 1 = c$ (However, for those expressions with an astrophysical application we have written the speed of light explicitely.). Greek indices are taken to run from 0 to 3 and Latin indices from 1 to 3; covariant derivatives are denoted with a semi-colon and partial derivatives with a comma. ",
        "conclusions": "\\label{conclusion} We have presented analytic general relativistic expressions for the electromagnetic fields internal and external to a slowly-rotating magnetized neutron star. The star is considered isolated and in vacuum, but no special assumption is made on the orientation of the dipolar magnetic field with respect to the rotation axis. The solutions to Maxwell equations have been considered both for an infinite and for a finite electrical conductivity.  In the first case, corresponding to stationary magnetic fields, we have shown that the general relativistic corrections due to the dragging of reference frames are not present in the form of the magnetic fields but emerge only in the form of the electric fields. In particular, we have shown that the frame-dragging provides an additional induced electric field which is analogous to the one introduced by the rotation of the star in the flat spacetime limit. In the case of finite electrical conductivity, on the other hand, corresponding to decaying magnetic fields, we have shown that corrections due both to the spacetime curvature and to the dragging of reference frames can be found in the induction equation. An interesting result obtained in this regime is that the rotation of the star eliminates the degeneracy in the components of the induction equation which remain therefore distinct. Furthermore, rotation and dipole misalignment do not eliminate the angular dependence in the induction equation and, as a result, an initially dipolar magnetic field might evolve towards a different configuration during its decay.  Because of their complexity, the evolution equations found for the magnetic field require a numerical integration which will discuss in detail in a forthcoming work (Rezzolla et al. 2000). There, we will also present direct comparisons between the flat and the curved spacetime solutions and quantify more precisely the importance of the general relativistic corrections.  One of the relevant aspects of the solutions presented in this paper is that they provide a lowest order analytic form for the electromagnetic field in the spacetime of a slowly rotating misaligned dipole subject to assumptions which, while giving simplifications, allow the major features of a realistic solution to be seen. In this sense, they reflect a rather general physical configuration and could therefore be used in a variety of astrophysical situations."
    },
    "0011/astro-ph0011498_arXiv.txt": {
        "abstract": "{ The EPIC focal plane imaging spectrometers on XMM-Newton use CCDs to record the images and spectra of celestial X-ray sources focused by the three X-ray mirrors. There is one camera at the focus of each mirror; two of the cameras contain seven MOS CCDs, while the third uses twelve PN CCDs, defining a circular field of view of 30$^{\\prime}$ diameter in each case. The CCDs were specially developed for EPIC, and combine high quality imaging with spectral resolution close to the Fano limit. A filter wheel carrying three kinds of X-ray transparent light blocking filter, a fully closed, and a fully open position, is fitted to each EPIC instrument. The CCDs are cooled passively and are under full closed loop thermal control. A radio-active source is fitted for internal calibration. Data are processed on-board to save telemetry by removing cosmic ray tracks, and generating X-ray event files; a variety of different instrument modes are available to increase the dynamic range of the instrument and to enable fast timing. The instruments were calibrated using laboratory X-ray beams, and synchrotron generated monochromatic X-ray beams before launch; in-orbit calibration makes use of a variety of celestial X-ray targets. The current calibration is better than 10\\% over the entire energy range of 0.2 to 10 keV. All three instruments survived launch and are performing nominally in orbit. In particular full field-of-view coverage is available, all electronic modes work, and the energy resolution is close to pre-launch values. Radiation damage is well within pre-launch predictions and does not yet impact on the energy resolution. The scientific results from EPIC amply fulfil pre-launch expectations. ",
        "introduction": "The EPIC instrument on XMM-Newton \\citep{jansen01} provides focal plane imaging and spectrometry for the three X-ray telescopes. Each telescope has an objective comprising a nested, 58 shell, Wolter 1 X-ray mirror \\citep{aschenbach01}, of focal length 7.5 metres, and geometric effective area 1500 cm$^{2}$; there is one EPIC at the focus of each telescope. Two of the telescopes are fitted with the X-ray gratings of the Reflection Grating Spectrometer \\citep{denherder01}. The gratings divert 50\\% of the flux out of the EPIC beams; with allowance for structural obscuration, 44\\% of the original flux reaches two of the EPIC cameras; these contain MOS CCDs \\citep{short98} and are referred to as the MOS cameras. The third telescope has an unobstructed beam; the EPIC instrument at the focus of this telescope uses PN CCDs \\citep{struder01} and is referred to as the PN camera. All three cameras have an identical forward section that contains a filter wheel, door, calibration source, radiation shielding, the interface to the spacecraft focal plane bulkhead, and the internal bulkhead that forms part of the camera vacuum enclosure. The rear part of each camera that contains the CCDs and the cooling system is different in construction for the MOS (figure~\\ref{mos}) and PN cameras. EPIC also includes the EPIC Radiation Monitor System, to record the ambient proton and electron flux \\citep{boer96}. It provides warning of a radiation flux increase to provide for automatic shut down of the instrument. This paper describes the common items and the MOS cameras, while an accompanying paper, by \\cite{struder01} describes the PN camera.  \\begin{figure} \\centering \\includegraphics[width=7cm]{Complete_MOS.ps} \\caption{The long conical radiators of the MOS enable the radiating surface to reach the plane of the spacecraft thermal shield to avoid parasitic heat-loads. } \\label{mos} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011384_arXiv.txt": {
        "abstract": "The redshift evolution of the galaxy cluster temperature function is a powerful probe of cosmology. However, its determination requires the measurement of redshifts for all clusters in a catalogue, which is likely to prove challenging for large catalogues expected from XMM--Newton, which may contain of order 2\\,000 clusters with measurable temperatures distributed around the sky. In this paper we study the {\\em apparent} cluster temperature, which can be obtained without cluster redshifts. We show that the apparent temperature function itself is of limited use in constraining cosmology, and so concentrate our focus on studying how apparent temperatures can be combined with other X-ray information to constrain the redshift. We also briefly study the circumstances in which non-thermal spectral features can give redshift information. ",
        "introduction": "Considerable attention has been devoted to the study of the evolution of the galaxy cluster temperature function with redshift, which promises to be an extremely powerful probe of the density parameter (Frenk et al.~1990; Oukbir \\& Blanchard 1992; Viana \\& Liddle 1996; Eke, Cole \\& Frenk 1996). Even very small numbers of high-mass, high-redshift clusters can rule out the critical-density paradigm; indeed, several authors claim that they have already done so (Henry 1997; Bahcall \\& Fan 1998; Eke et al.~1998) though this remains controversial (Sadat, Blanchard \\& Oukbir 1998; Reichart et al.~1999; Viana \\& Liddle 1999). In order to fully apply this method, the cluster masses must be accurately determined and the usual technique is to use the gas temperature as measured from the X-ray emission. In addition the cluster redshift is required, both to place it correctly in the evolutionary sequence, and because the redshift is needed to convert the apparent temperature into the actual cluster temperature.  For existing catalogues of clusters for which the temperatures could be estimated, obtaining the redshifts proved a manageable task, as the number of clusters with sufficient photon counts to allow temperature determination was small. This is set to change with observations by the XMM--Newton (hereafter just XMM) satellite; in a recent paper (Romer et al.~1999) we showed that a planned serendipitous cluster survey which will analyze all XMM--EPIC frames suitable for serendipitous cluster detection, {\\sc Xcs}\\footnote{See {\\tt www.xcs-home.org}\\ \\ for further details.}, may contain as many as 10\\,000 galaxy clusters of temperature 2 keV and above, of which around 2\\,000 may have sufficient photon counts to allow the temperatures to be accurately estimated without further X-ray observations. Given that these will be distributed more or less randomly across the sky, follow-up to obtain spectroscopic redshifts represents a substantial task.  The main focus of this paper is on the use of the full available X-ray information to optimize the follow-up efficiency onto the high-redshift population.  Although a survey like the {\\sc Xcs} is likely to contain around two thousand clusters with high enough photon counts to permit an accurate temperature estimate, the thermal bremsstrahlung spectrum only gives the apparent temperature \\begin{equation} T_{{\\rm app}} = \\frac{T}{1+z} \\,, \\end{equation} with the true temperature not being known until the redshift is determined. For very luminous clusters this degeneracy may be broken by visible spectral lines such as the Iron K line complex at $7$ keV, but this will be challenging for most clusters and we defer discussion of this possibility until the end of the paper.  In this paper we discuss several aspects of apparent temperatures and the estimation of cluster redshifts from X-ray data. Apparent temperatures of clusters were first discussed by Oukbir \\& Blanchard (1997) in the context of the ROSAT all-sky survey. They also noted the curious point that even in the absence of redshifts the apparent temperature should still be a good estimator of relative cluster masses; for example in a critical-density model scaling laws predict $M \\propto T^{3/2}/(1+z)^{3/2} \\propto T_{{\\rm app}}^{3/2}$.  We focus on issues of apparent temperatures relevant to the XMM satellite. First of all, we analyze whether the apparent temperature function (that is, the number density of clusters observed above a given apparent temperature $T_{{\\rm app}}$) might in itself prove a useful probe of cosmology.  The answer will be that it proves of limited use, demonstrating the importance of determining cluster redshifts at the earliest possible stage.  In that light, we go on to consider how other X-ray observables can be combined with the apparent temperatures in order to constrain the redshifts, particularly with a view to eliminating low-redshift clusters from the follow-up candidate list.  The main observables are the angular size and the apparent luminosity of the clusters, and it is the latter which proves powerful in combination with the apparent temperature.  We also briefly study {\\sc xspec} spectral simulations to assess the likelihood of redshift determination from X-ray spectral lines. ",
        "conclusions": "In this paper we have studied various aspects of apparent cluster temperatures relevant to XMM observations. The apparent temperature function will be readily derived from XMM data, but will primarily be useful in constraining the degenerate combination familiar from low-redshift cluster number density studies, $\\sigma_8 \\Omega_0^{-\\alpha}$ with $\\alpha \\sim 0.5$, and is unlikely to usefully probe the density parameter itself.  To fully exploit a galaxy cluster catalogue, cluster redshifts are essential, and we have studied how other X-ray properties can be combined with the apparent temperature to select follow-up candidates efficiently. Assuming point source contamination can be recognized from the high-resolution imaging, it appears that the cluster flux, when combined with the apparent temperature, will yield a good indication of the cluster redshift, especially once early XMM observations have been used to calibrate the relation.  Finally, we remark that further useful information towards estimating the redshifts may come from imaging data in the optical and infra-red once cluster candidates have been identified in the X-ray. A good example is the K band luminosity of the brightest cluster galaxy, which exhibits a tight correlation with redshift for X-ray luminous clusters, as shown by Collins \\& Mann (1998) and Burke, Collins \\& Mann (2000). In many cases it may also be possible to carry out photometric redshifting of cluster galaxies, for example using Sloan Digital Sky Survey and VISTA data."
    },
    "0011/hep-ex0011001_arXiv.txt": {
        "abstract": "We present the first direct experimental evidence for the charge excess in high energy particle showers predicted nearly 40 years ago by Askaryan. We directed bremsstrahlung photons from  picosecond pulses of 28.5 GeV electrons at the SLAC Final Focus Test Beam facility into a 3.5 ton silica sand target, producing electromagnetic showers several meters long. A series of antennas spanning 0.3 to 6 GHz were used to detect strong, sub-nanosecond radio frequency pulses produced whenever a shower was present. The measured electric field strengths are consistent with a completely coherent radiation process. The pulses show 100\\% linear polarization, consistent with the expectations of Cherenkov radiation. The field strength versus depth closely follows the expected particle number density profile of the cascade, consistent with emission from excess charge distributed along the shower. These measurements therefore provide strong support for experiments designed to detect high energy cosmic rays and neutrinos via coherent radio emission from their cascades. ",
        "introduction": "During the development of a high-energy electromagnetic cascade in normal matter, photon and electron scattering processes pull electrons from the surrounding material into the shower.  In addition, positrons in the shower annihilate in flight.  The combination of these processes should lead to a net 20-30\\% negative charge excess for the comoving compact body of particles that carry most of the shower energy. G. A. Askaryan\\cite{Ask62} first described this effect, and noted that it should lead to strong coherent radio and microwave Cherenkov emission for showers that propagate within a dielectric. The range of wavelengths over which coherence obtains depends on the form factor of the shower bunch---wavelengths shorter than the bunch length suffer from destructive interference and coherence is lost. However, in the fully coherent regime the radiated energy scales quadratically with the net charge of the particle bunch, and at ultra high energies the resulting coherent radio emission may carry off a significant fraction of the total energy in the cascade.  The plausibility of Askaryan's arguments combined with more recent modeling and analysis\\cite{Mark86,ZHS92,Alv96,Alv97,Alv98} has led to a number of experimental searches for high energy neutrinos by exploiting the effect at energies from $\\sim 10^{16}$ eV in Antarctic ice\\cite{fri96,Bes99} up to $10^{20}$ eV or more in the lunar regolith, using large ground-based radio telescopes\\cite{Zhe88,Dag89,Han96,Gor99}. Radio frequency pulses have been observed for many years from extensive air showers\\cite{Jel66,Feg68}. However, it has been shown\\cite{KL66,All71} that the most likely source of this emission is a form of Lorentz-boosted dipole radiation from geomagnetic charge separation in the air shower, rather than the Askaryan effect. Thus neither the charge asymmetry nor the resulting coherent Cherenkov radiation has ever been observed.  In a previous paper\\cite{P1} we have described initial efforts to measure the coherent radio-frequency (RF) emission from electron bunches interacting in a solid dielectric target consisting of 360~kg of silica sand.  That study, done with relatively low-energy electrons (15 MeV), demonstrated the presence of coherent radiation in the form of extremely short and intense microwave pulses detectable over a wide frequency range. These results, while useful for understanding the coherent RF emission processes from relativistic charged particles, could not directly test the development of a shower charge excess, as Askaryan predicted. Also, because the particles were charged, passage of the beam through any interface induced strong RF transition radiation (TR), which obscured the presence of Cherenkov radiation (CR).  We report here on measurements made at the Stanford Linear Accelerator Center (SLAC) in which the use of high-energy photons, rather than low-energy electrons, has enabled us to clearly observe microwave Cherenkov radiation from the Askaryan effect. The electromagnetic showers thus produced in our target yielded strong, coherent, sub-ns RF pulses, consistent in every way observed with the predictions. As we will show here, our results conclusively support the reported sensitivity of the experiments noted above. ",
        "conclusions": "We have demonstrated that the radiation we have observed is coherent, and 100\\% linearly polarized. The plane of polarization coincides with the plane containing the antenna and shower axis. The radiation is pulsed with time durations much shorter than the inverse bandwidth of our antennas. The strength of the pulse is strongly correlated to the size (in particle number) of the shower region that appears to produce it. All of these observed characteristics are consistent with the hypothesis that what we have observed is coherent Cherenkov radiation. Because of the strong correlation with the shower profile, and the physical constraint that a shower with no net excess charge cannot radiate, we conclude that excess charge production along the shower is the source of the propagating charge in the silica dielectric which leads to the CR observed. These conclusions are strengthened by the fact that the pulse strength is consistent with Monte Carlo predictions for such showers.  Our observations are inconsistent with radiation from geomagnetic charge separation as observed in extensive air showers. The most striking evidence for this is the fact that the plane of polarization is clearly aligned with the shower axis rather than the local geomagnetic dip angle ($62^{\\circ}$). Given the approximate east-west orientation of the shower axis, boosted dipole radiation from geomagnetic charge separation should produce an electric field polarization with significant components orthogonal to what we have observed.  We note that in all cases our measurements are likely to be made in near-field conditions, and we have not attempted to correct for these effects. Recent studies\\cite{Bun00,Alv00} have begun to treat these issues, but are not straightforward to apply in our case. In any case, near field effects should generally {\\em decrease} the measured field strengths relative to far field measurements.  The total energy of our cascades is as high as $10^{19}$ eV, but these showers consist of the superposition of many lower-energy showers. Higher-energy effects\\cite{Alv98,Alv99} that elongate the shower during its development are not present. These effects all tend to increase the total tracklength of the shower, at the expense of a lower instantaneous particle number density. The net effect is that the total radiated power is expected to be approximately conserved, but the angular spectrum can change significantly, with a predicted sharpening of the Cherenkov angular distribution for high energy showers. We have not measured the angular spectrum of the radio pulses, and thus any scaling of our results to high energy showers should consider corrections for these effects, which are likely to increase the peak field strength in far-field measurements.  Since the field strength scales linearly in shower energy and inversely with distance from the source, extrapolations to determine the energy threshold of existing experiments is straightforward, after correcting for the differences in material properties. For Antarctic ice experiments, the use of the existing simulations \\cite{ZHS92,Alv98,fri96,Alv00} appears completely justified. For experiments observing the lunar regolith\\cite{Zhe88,Dag89,Han96,Gor99}, silica sand shares many similarities with the lunar surface material, and the expected cascade energy threshold, by direct scaling of our results, is $\\sim 2 \\times 10^{19}$ eV, and may be somewhat lower depending on the angular intensity effects discussed above.  We conclude that, in combination with our previous measurements of coherent RF transition radiation~\\cite{P1}, these results have established a firm experimental basis for radio-frequency detection of high energy cascades in solid media, either through interaction within a dielectric (for CR), or via passage through dielectric interfaces (for TR). Above cascade energies of $\\sim 10^{16}$~eV, these secondary emission processes become dominant over others (for example, optical Cherenkov or fluorescence emission) in the number of quanta produced\\cite{Mark86}. Thus experiments designed to exploit this effect in the detection of ultra-high energy particles can now be pursued with even greater confidence."
    },
    "0011/astro-ph0011047_arXiv.txt": {
        "abstract": "We have obtained HST-NICMOS observations of five of M31's most metal rich globular clusters: G1, G170, G174, G177 \\& G280.  For the two clusters farthest from the nucleus we statistically subtract the field population and estimate metallicities using $K$-$(J-K)$ color-magnitude diagrams (CMDs).  Based on the slopes of their infrared giant branches we estimate [Fe/H] $=-1.22\\pm0.43$ for G1 and $-0.15\\pm0.37$ for G280. We combine our infrared observations of G1 with two epochs of optical HST-WFPC2 $V$-band data and identify at least one LPV based on color and variability.  The location of G1's giant branch in the $K$-$(V-K)$ CMD is very similar to that of M107, indicating a higher metallicity than our purely infrared CMD: [Fe/H]$\\sim -0.9\\pm0.2$.  For the three central clusters, which are too compact for accurate cluster star measurements, we present integrated cluster magnitudes and field CMDs.  The $K$-band luminosity functions (LFs) of the upper few magnitudes of G1 and G280, as well as for the fields surrounding all clusters, are indistinguishable from the LF measured in the bulge of our Galaxy.  This indicates that these clusters are very similar to Galactic clusters, and at least in the surrounding fields observed, there are no significant populations of young luminous stars.  For the field surrounding G280, we estimate the metallicity to be $-1.3$ from the slope of the giant branch, with a spread of $\\sigma_{[Fe/H]} \\sim 0.5$ from the width of the giant branch.  Based on the numbers and luminosities of the brightest giants, we conclude that only a small fraction of the stars in this field could be as young as 2 Gyr, while the majority have ages closer to 10 Gyr. ",
        "introduction": "\\label{sec:introduction}  Globular clusters (GC) occupy a very special position in modern astrophysics.  They provide the most stringent tests of stellar evolution theory, represent ideal templates for stellar population synthesis studies, allow age dating of galaxies with unrivaled precision, etc.  The GC families of the Milky Way and its satellite galaxies, the Magellanic Clouds and the Fornax dSph galaxy, have been thoroughly studied from both the ground and space, especially with HST. The next nearest major GC family belongs to M31.  To study individual stars in GC systems much more distant than M31 will require the next generation of ground or space based telescopes.  The GC systems of the the Milky Way and M31 appear to be quite similar in luminosity and color ranges \\citep{BBBF1987, EW1988, HBK1991, FPC1980b, DGMC1997, SBA1996}.  However, some of M31's bulge globulars appear to have line strengths as strong as those of giant ellipticals, suggesting metallicities considerably greater than any known Galactic globular \\citep{BJSA1992, JAB1992, Jab1997}. \\citet{BFGK1984} has also noted that M31's GCs appear to follow different H$\\beta$ and CN correlations with Mg$_2$ with respect to Galactic globulars; they attributed this to the whole M31 GC family being systematically younger. Renzini (1986) pointed out that other interpretations were also possible, such as the chemical enrichment history of the two spheroids having proceeded with slightly different time scales, resulting in different element ratios.  HST observations of M31's GCs in the optical (with the goal of stellar photometry) started soon after the first refurbishing mission, with both FOC and WFPC2 \\citep{FBCC1996, RMFN1996, AGLB1996}, and now include some of the most metal rich clusters \\citep{JCMS2000}.  The study of metal rich M31 globulars complements similar studies of Galactic bulge globulars within this still poorly known, yet crucial part of age-metallicity space.  Near-IR observations are essential to study the brightest giants in an old, metal rich population, especially to determine bolometric luminosities.  Metal rich globular and bulge giants, which are the brightest bolometrically and in the near-IR, are in contrast, many magnitudes fainter at optical wavelengths due to severe molecular blanketing.  In extreme cases, $(V-K)$ can be as great as $\\sim 10$ at the top of the AGB \\citep{FW1987, GOMR1998} and these stars are likely to have escaped detection in M31 even with WFPC2.  Since the main sequences of the old populations in M31 are currently out of reach, one is forced to appeal to an alternative age indicator. Theory predicts that the highest luminosity reached on the AGB is a function of age \\citep{IR1983}, a prediction extensively verified by observations of clusters in the Magellanic Clouds \\citep{MA1986, FMB1990}.  In metal poor Galactic globulars, no stars are brighter than the theoretical RGB tip, as expected for a $\\sim 14$ Gyr population. However, the AGB of more metal rich clusters ([Fe/H]$>-1$) extends $\\sim 1$ magnitude above the RGB \\citep{FPC1983, FE1988, GRO1997}, as it does in the Galactic bulge \\citep{FW1987}.  For clusters, though, these bright stars are all LPVs.  For both clusters and bulge population these luminous stars are now generally ascribed to high metallicity, rather than to young age \\citep{FW1987, GRO1997}.  Hence, the presence of stars brighter than the RGB tip does not guarantee an intermediate age; consideration of their color, luminosity {\\it and} frequency in the parent population is required before drawing conclusions about ages \\citep{Ren1993}.  In Cycle 7 we proposed to obtain NICMOS $JHK$ images of 5 metal rich globular clusters in M31.  From these observations, we planned to achieve the following scientific goals: (1) explore the upper end of the GC luminosity function; (2) derive independent estimates for the metallicity of the M31 clusters and their adjacent fields from the slope and location of the NIR RGB \\citep{KFTP1995}; (3) determine the frequency of luminous RGB and AGB stars (including LPVs) per unit luminosity; (4) compare the cluster results with observations of M31 field stars (5) compare the properties of the luminous stars in M31's metal rich clusters to those in the Galactic bulge and Galactic globulars; (6) integrate these results with optical photometry and spectroscopy and explore implications for stellar evolution theory and the interpretation of the integrated light of distant galaxies.  One critical issue deserves special attention if any of these goals are to be attained: the effect of stellar crowding.  We \\citep[][hereafter Paper I]{SFFJ2001}, have carefully analyzed the effects of blending on our NICMOS data.  Through the creation of completely artificial clusters, we have calculated threshold- and critical-blending limits for each cluster and surrounding field.  These limits determine the proximity to each cluster where reliable photometry can be obtained. These simulations allow us to quantify and correct for the effects of blending on the GB slope and width at different surface brightness levels.  This paper is organized as follows.  Section \\ref{sec:observations} presents the reasons for selecting each cluster, and the details of the observations.  Section \\ref{sec:data_reduction} describes the reduction procedures, and gives a brief summary of the procedures and results on blending from Paper I.  Section \\ref{sec:photometry} presents the integrated photometry of the clusters, the CMDs and luminosity functions of the clusters and their surrounding fields, and metallicity estimates for G1, G280, and the G280 field.  Our conclusions are summarized in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We first present surface brightness profiles of all clusters, and with the blending analysis presented in Paper I, determine radial photometric limits for each cluster.  We then give integrated photometry of all clusters except G1, which was not fully on the detector.  For the G1 cluster, we present the infrared CMD, and estimate the metallicity as [Fe/H]$= -1.22\\pm0.43$ from the slope of the giant branch.  Based on the width of the giant branch, which shows no significant spread in color over what is expected from measurement errors alone $(\\sigma_{(J-H)}=0.06)$, we conclude that there is no significant metallicity spread in the cluster.  We combine our infrared observations of G1 with two epochs of optical $V$-band HST-WFPC2 data, revealing that several of the brightest stars in the cluster are LPVs. The shape and position of the GB in the $K$-$(V-K)$ CMD are similar to that of M107, indicating a metallicity of [Fe/H]$=-0.9\\pm0.2$.  However since the infrared GB slope technique uses such a small range in luminosity, we place more weight on the higher value from the optical-infrared CMD.  For the G280 observations, we divide the frame into cluster and field at $5''$ from the cluster center.  We statistically subtract the field population from the cluster, and present both the cluster and field CMDs.  Fitting the giant branch, we find a cluster metallicity of [Fe/H]$= -0.15\\pm0.37$.  As in G1, we see no evidence for a metallicity spread in the cluster based on the width of the GB $(\\sigma_{(J-H)}=0.07)$.  Fitting the GB of the G280 field, we find a metallicity of $-1.3 \\pm 0.6$.  The large error on the metallicity is indicative of the large color spread, which we estimate to be $\\sigma_{[Fe/H]} \\sim 0.5$ dex from the width of the GB.  This is not surprising, since this field has contributions from both the disk (85\\%) and bulge (15\\%).  What is surprising is that, in this field which is 85\\% disk, we see no obviously bright, young stars.  Using the brightest star in the field as the tip of the AGB, at $M_{bol}=-5$, we estimate an age of $\\sim 5$ Gyr. However, if the disk component were this young, we would expect to see $\\sim 30$ stars brighter than $M_{bol}=-4$, but we see only 7.  A more likely scenario is that there are just a few young disk stars in the field, while the majority of the disk population is closer to $\\sim 10-15$ Gyr, thus lowering the AGB tip to $M_{bol} \\sim -4.5$.  The three central clusters, G170, G174 \\& G177, are all too compact to extract cluster star photometry.  We thus present the field CMDs and luminosity functions without trying to separate the cluster and field star contributions.  The surface brightness of the G170 field is within acceptable limits, so we perform a linear fit to the GB, but do not try to estimate the metallicity, as the blending correction would be uncomfortably large.  The G174 and G177 fields, on the other hand, are both above the threshold-blending surface brightness limit.  This implies that measurements of all but the brightest stars in these fields are potentially affected by blending, so we refrain from even fitting their GBs.  Finally we presented the cluster (G1 and G280) and field luminosity functions with the LF measured in Baade's Window.  The luminosity functions of G1 and G280 both have a sharp bright-end cutoff at $M_K \\sim -6.5$, consistent with observations of Galactic globulars.  The fields surrounding the clusters have LFs which are indistinguishable from that measured in the Galactic bulge.  Thus, at least in the fields observed, there is no significant population of young luminous stars in the bulge of M31.  \\"
    },
    "0011/astro-ph0011271_arXiv.txt": {
        "abstract": "{ We report on the first deep X-ray survey with the {\\em XMM-Newton} observatory during the performance verification phase. The field of the Lockman Hole, one of the best studied sky areas over a very wide range of wavelengths, has been observed. A total of $\\sim$ 100 ksec good exposure time has been accumulated. Combining the images of the {\\em European Photon Imaging Camera} (EPIC) detectors we reach a flux limit of 0.31, 1.4 and $2.4 \\times 10^{-15}~{\\rm erg}~{\\rm cm}^{-2}~{\\rm s}^{-1}$, respectively in the 0.5-2, 2-10, and 5-10 keV band. Within an off-axis angle of 10 arcmin we detect 148, 112 and 61 sources, respectively. The log(N)-log(S) relation in the three bands is compared with previous results. In particular in the 5-10 keV band these observations present the deepest X-ray survey ever, about a factor 20 more sensitive than the previous {\\em BeppoSAX} observations. Using X-ray spectral diagnostics and the set of previously known, spectroscopically identified {\\em ROSAT} sources in the field, the new sources can be classified. {\\em XMM-Newton} detects a significant number ($\\sim$ 40\\%) of X-ray sources with hard, probably intrinsically absorbed X-ray spectra, confirming a prediction of the population synthesis models for the X-ray background. ",
        "introduction": "\\label{sec:intr}  Deep X-ray surveys indicate that the cosmic X-ray background (XRB) is largely due to accretion onto supermassive black holes, integrated over cosmic time. In the soft (0.5-2 keV) band 80-90\\% of the XRB flux has been resolved using {\\em ROSAT} and recent {\\em Chandra} surveys (Hasinger et al. \\cite{hasi98a}, Mushotzky et al. \\cite{mushotzky00}, Giacconi et al. \\cite{giacconi00}). In the harder (2-10 keV) band 25-30\\% of the background have been resolved in {\\em ASCA} and {\\em BeppoSAX} surveys (Ueda et al. \\cite{ueda98}, Cagnoni et al.  \\cite{cagnoni98}, Giommi et al. \\cite{giommi00}), and more than 60\\%, when the recent {\\em Chandra} surveys are included. Surveys in the very hard (5-10 keV) band have been pioneered using {\\em BeppoSAX} and resolve about 30\\% of the XRB (Fiore et al. \\cite{fiore99}). Those X-ray surveys with a high degree of completeness in optical spectroscopy find predominantly Active Galactic Nuclei (AGN) as counterparts of the faint X-ray source population (Bower et al. \\cite{bower96}, Schmidt et al. \\cite{schmidt98}, Zamorani et al. \\cite{zamorani99}, Akiyama et al. \\cite{akiyama00}), mainly X-ray and optically unobscured AGN (type-1 Seyferts and QSOs) but also a smaller fraction of obscured AGN (type-2 Seyferts). Spectroscopic identifications of the {\\em BeppoSAX} and {\\em Chandra} surveys are still far from complete, however a mixture of obscured and unobscured AGN seems to be the dominant population in these samples, too (Fiore et al. \\cite{fiore00}, Barger et al. \\cite{barger00}, Giacconi et al. \\cite{giacconi00}). The most recent AGN X-ray luminosity function, derived from the {\\em ROSAT} surveys, shows evidence for luminosity-dependent density evolution and indicates a constant QSO space density at redshifts $2 < z < 4$ (Hasinger \\cite{hasi98b}, Miyaji et al. \\cite{miyaji00}) unlike optical QSO luminosity functions (Schmidt et al. \\cite{schmidt95}, Fan et al. \\cite{fan00}).  The X-ray observations are consistent with population synthesis models based on unified AGN schemes (Setti \\& Woltjer \\cite{setti89}, Madau et al. \\cite{madau94}, Comastri et al. \\cite{comastri95}, Gilli et al. \\cite{gilli99}), which explain the hard spectrum of the X-ray background by a mixture of absorbed and unabsorbed AGN, folded with the corresponding luminosity function and its cosmological evolution. According to these models most AGN spectra are heavily absorbed and about 80\\% of the light produced by accretion will be absorbed by gas and dust (Fabian et al. \\cite{fabian98}). However, these models are not unique and contain a number of hidden parameters, so that their predictive power remains limited (e.g. Hasinger \\cite{hasi00}). In particular they require a substantial contribution of high-luminosity obscured X-ray sources (type-2 QSOs), which so far have not been detected in sufficient quantities (see the discussion in Halpern et al. \\cite{halpern99}). The large throughput and the unprecedented hard X-ray sensitivity of the telescopes aboard the recently launched {\\em XMM-Newton} observatory (hereafter {\\em XMM}; Jansen et al. \\cite{jansen01}) will ultimately yield spectra of the faint X-ray sources and constrain the evolution of their physical properties, in particular the X-ray absorption.  Here we present results of the first deep survey taken with {\\em XMM} in the {\\em Lockman Hole}, one of the best studied sky areas at all wavelengths. This paper concentrates on the X-ray data analysis. We show combined images from the EPIC pn-CCD (Str\\\"uder et al. \\cite{strueder01}) and MOS CCD cameras (Turner et al. \\cite{turner01}) and the derived source counts in different energy bands. With the help of X-ray colour-colour diagrams and the previously identified sources in this field we show that it is possible to obtain a coarse source classification based on {\\em XMM} data alone. ",
        "conclusions": "\\label{sec:disc}  We have shown the first log(N)-log(S) relations based on the {\\em XMM-Newton} observatory. Given the still existing systematic uncertainties, the data is fully consistent with the {\\em ROSAT} and {\\em Chandra} source counts in the 0.5-2 keV band. This demonstrates on one hand that the combination of EPIC detectors is not yet confusion limited in a 100 ksec observation, on the other hand that cosmic variance between different fields does not affect the source counts significantly at the currently achieved flux levels, at least not in the soft band.  In the 2-10 keV band there is an inconsistency of about 40\\% between the two recent {\\em Chandra} datasets by Mushotzky et al. \\cite{mushotzky00} and Giacconi et al. \\cite{giacconi00}, the latter one having a lower normalisation. The new {\\em XMM} data are consistent with the Giacconi et al. log(N)-log(S), maybe even somewhat flatter, and clearly confirm a break in the slope compared to the quasi-Euclidean behaviour at brighter fluxes. In the 5-10 keV band the {\\em XMM} data go more than an order of magnitude deeper than the previous {\\em BeppoSAX} counts (Fiore et al. \\cite{fiore99}). There is so far relatively little deviation from a Euclidean slope and the data is fully consistent with the predictions from recent population synthesis models for the X-ray background (Gilli et al. \\cite{gilli99}). Adding up the source counts, we resolve about 60\\% of the 5-10 keV X-ray background.  The diagnostic power of {\\em XMM} lies in its wide energy band and its unprecedented sensitivity in the hard band. With the help of X-ray colour-colour diagrams and the ``training set'' of about 60 previously identified {\\em ROSAT} sources in the same field it is possible to characterise the new XMM sources as typically harder, probably intrinsically absorbed sources. A small number of objects with similar X-ray colours has already been identified in the deepest ROSAT survey (Lehmann et al. \\cite{lehmann00}), they are type-2 Seyferts or unidentified objects with extremely red optical/NIR colours ($R-K > 5$). The new XMM source population is therefore very likely dominated by obscured AGN, as predicted by the AGN population synthesis models for the X-ray background."
    },
    "0011/astro-ph0011101_arXiv.txt": {
        "abstract": "The dynamic orchestration of starbirth activity in the starburst-ringed galaxy M94 (NGC 4736) is investigated using images from the Ultraviolet Imaging Telescope (FUV-band), Hubble Space Telescope (NUV-band), Kitt Peak 0.9-m telescope (H$\\alpha$, R, and I bands), and Palomar 5-m telescope (B-band), along with spectra from the International Ultraviolet Explorer and Lick 1-m telescopes. The wide-field UIT image shows FUV emission from (a) an elongated nucleus, (b) a diffuse inner disk, where H$\\alpha$ is observed in {\\it absorption}, (c) a bright inner ring of H II regions at the perimeter of the inner disk (R = 48$''$ = 1.1 kpc), and (d) two 500-pc size knots of hot stars exterior to the ring on diametrically opposite sides of the nucleus (R = 130$''$ = 2.9 kpc). The HST/FOC image resolves the NUV emission from the nuclear region into a bright core and a faint 20$''$-long ``mini-bar'' at a position angle of 30 deg. Optical and IUE spectroscopy of the nucleus and diffuse inner disk indicates a $\\sim$10$^{7-8}$ yr-old stellar population from low-level starbirth activity blended with some LINER activity. Analysis of the H$\\alpha$, FUV, NUV, B, R, and I-band emission along with other observed tracers of stars and gas in M94 indicates that most of the star formation is being orchestrated via ring-bar dynamics involving the nuclear mini-bar, inner ring, oval disk, and outer ring.  The inner starburst ring and bi-symmetric knots at intermediate radius, in particular, argue for bar-mediated resonances as the primary drivers of evolution in M94 at the present epoch. Similar processes may be governing the evolution of the ``core-dominated'' galaxies that have been observed at high redshift.  The gravitationally-lensed ``Pretzel Galaxy'' (0024+1654) at a redshift of $\\sim$1.5 provides an important precedent in this regard. ",
        "introduction": "Star-forming rings or ``pseudorings'' are common to early and intermediate-type spiral galaxies (cf.\\markcite{Athan85} Athanassoula \\& Bosma 1985; \\markcite{buta95b}Buta, Purcell, \\& Crocker 1995; \\markcite{bc96}Buta \\& Combes 1996), including our own Milky Way galaxy (cf. \\markcite{gm82}Gusten \\& Mezger 1982; \\markcite{Clemens88}Clemens, Sanders, \\& Scoville 1988; \\markcite{Waller90a}Waller 1990a). Such ring-like accumulations of gas and associated starbirth activity may have helped to build the inner parts of many primeval disk galaxies (\\markcite{Friedli95}Friedli \\& Benz 1995), as exemplified by the recent discovery of a starburst-ringed galaxy at $z \\sim 1.5$ (\\markcite{Colley96}Colley, Tyson, \\& Turner 1996; \\markcite{Tyson97}Tyson et al. 1997).  The formation and maintenance of these starburst rings are often attributed to orbital resonances with rotating bar or ``oval'' asymmetries in the stellar disks (cf. \\markcite{Combes94}Combes 1994, \\markcite{Byrd94}Byrd et al. 1994; \\markcite{Combes95}Combes et al. 1995; \\markcite{bc96}Buta \\& Combes 1996 and references therein). However, other dynamical mechanisms --- including gravitational instabilities (\\markcite{Elmegreen92}Elmegreen 1992, \\markcite{Elmegreen94}1994; \\markcite{Kenney97}Kenney \\& Jogee 1997), outward propagating star formation (\\markcite{Walker88}Walker, Lebofsky, \\& Rieke 1988; \\markcite{Waller92}Waller, Gurwell, \\& Tamura 1992), and even radially-driven pileups from nuclear outbursts (Waller et al. 1992; \\markcite{Tenorio97}Tenorio-Tagle et al. 1997) --- may play significant roles in orchestrating some of the starburst rings that are observed. \\footnote{Collisionally-induced ``ring galaxies'' such as the Cartwheel Galaxy are thought to be morphologically and dynamically distinct from the more common ``ringed galaxies'' considered herein (cf. \\markcite{Athan85}Athanassoula \\& Bosma 1985; \\markcite{Marcum92} Marcum et al. 1992).}  As the closest early-type spiral galaxy of low inclination, M94 (NGC 4736) has received concentrated attention from both observers and theorists.  This (R)SA(r)ab-type galaxy (\\markcite{rc3}de Vaucouleurs et al. 1991) is noted for its inner ring of ongoing starburst activity (R $\\approx$ 45$''$), oval stellar distribution at intermediate radius (R $\\approx$ 220$''$) (cf. \\markcite{Mulder93}Mulder \\& van Driel 1993; \\markcite{Mulder95}Mulder 1995; \\markcite{Mollenhoff95} Mollenhoff, Matthias, \\& Gerhard 1995), and outer stellar ring near its de Vaucouleurs radius (R$_{25}$ $\\approx$ 330$''$). {\\bf Figure 1} (extracted from the Digital Sky Survey using the SkyView utility [McGlynn, Scollick, \\& White 1996])\\footnote{NASA's SkyView facility (http://skyview.gsfc.nasa.gov) was developed and is maintained under NASA ADP Grant NAS5-32068 at NASA's Goddard Space Flight Center.} shows the outermost portions of M94, highlighting the oval disk and outer pseudoring.  The inner starburst ring is a prominent source of H$\\alpha$, H I, and CO emission (\\markcite{smith91}Smith et al. 1991; \\markcite{mulder93}Mulder \\& van Driel 1993; \\markcite{Gerin91} Gerin, Casoli, \\& Combes 1991). The discovery of compact thermal \\& nonthermal radio sources in the ring (\\markcite{Duric88}Duric \\& Dittmar 1988) indicates the presence of dense H II regions and young SNRs. The ring's velocity field can be described by a combination of circular rotation with velocities of order 200 km/s and residual non-circular motions of order 15 km/s (\\markcite{Mulder95}Mulder 1995) to 25 km/s (\\markcite{buta88}Buta 1988), depending on the adopted inclination and major axis position angle.  Interior to the ring, the bright bulge and inner disk show twisted isophotes at red and near-IR wavelengths, indicative of a weak bar-like distortion (\\markcite{Beckman91}Beckman et al. 1991; \\markcite{Shaw93}Shaw et al. 1993; \\markcite{Mollenhoff95} Mollenhoff et al. 1995). FIR and CO observations interior to the ring reveal a rich ISM with gas surface densities exceeding that of the ring (\\markcite{Smith94}Smith \\& Harvey 1994; \\markcite{Garman86}Garman \\& Young 1986; \\markcite{Gerin91}Gerin et al. 1991; \\markcite{Wong00}Wong \\& Blitz 2000).  Optical spectroscopy of the nuclear region yields LINER-type emission lines along with absorption lines from the circumnuclear stellar population, consistent with an early main-sequence stellar turnoff (A4--A7) and corresponding age of $\\sim$500 Myr (\\markcite{Pritchett77} Pritchett 1977; \\markcite{Keel83}Keel 1983; \\markcite{Taniguchi96}Taniguchi et al. 1996). Further support for a young central population comes from NIR spectroscopy which shows deep CO absorption bands from red giant and asymptotic giant branch stars of similar age (\\markcite{Walker88}Walker et al. 1988).  These authors have proposed an outward propagating mode of star formation, whereby NGC 253, M82, M94, and M31 represent increasingly evolved versions of the same starbursting sequence. Although the kinematics of the ring show very little evidence for outward {\\it expanding motions} (contrary to prior claims of bulk expansion [\\markcite{kruit74}van der Kruit 1974; \\markcite{kruit76}1976]), they also do not preclude a scenario for {\\it radially propagating} star formation. Other investigators have modeled the inner and outer rings in terms of {\\it resonant dynamics} mediated by bar or ``oval'' potentials interior to the rings (\\markcite{Gerin91}Gerin et al. 1991; \\markcite{Shaw93} Shaw et al. 1993; \\markcite{Mollenhoff95}Mollenhoff et al. 1995; \\markcite{Mulder96}Mulder \\& Combes 1996), with the observed non-circular motions resulting from dispersion orbits near the Lindblad resonances (\\markcite{buta88}Buta 1988).  In this paper, we present and discuss new observational clues to the dynamical mechanisms governing the star formation in M94. The ultraviolet images obtained by the Ultraviolet Imaging Telescope, in particular, reveal hitherto unrecognized patterns of recent star formation whose presence lends further support to the hypothesis of galaxy evolution via bar-mediated resonances. The various imaging and spectroscopic observations and reductions are described in Section 2.  The resulting FUV, NUV, H$\\alpha$, R, and I-band emission morphologies are presented and compared in Section 3.  Radial intensity profiles and other photometric comparisons are discussed in Section 4.  UV and optical spectroscopy of the inner disk and nucleus is presented in Section 5.  Kinematic properties and inferred dynamical scenarios are considered in Section 6.  Our summary of the wavelength-dependent morphological, spectro-photometric, and dynamical properties of M94 appears in Section 7, wherein evolutionary implications are discussed.  In the following Sections, we assume a distance to M94 of 4.6 (75/H$_{\\circ}$) Mpc, based on the galaxy's recession velocity of 345 km/s with respect to the Local Group (\\markcite{Sandage81} Sandage \\& Tammann 1981).  The corresponding spatial scale is 22.3 pc/arcsec.  Unresolved sources imaged by the HST/FOC provide additional constraints on the distance, as discussed in Section 4.  We adopt a nominal inclination of 40$^{\\circ}$ and major-axis position angle of 120$^{\\circ}$, while recognizing that both of these quantities may vary significantly with radius and with measuring technique (e.g. morphological vs. kinematic determinations) (\\markcite{Bosma77}Bosma, van der Hulst, \\& Sullivan 1977; \\markcite{buta88}Buta 1988; \\markcite{Mulder93}Mulder \\& van Driel 1993; \\markcite{Mulder95}Mulder 1995; \\markcite{Mollenhoff95}Mollenhoff et al. 1995; \\markcite{Wong00}Wong \\& Blitz 2000). ",
        "conclusions": "Through UV-Optical imaging and spectroscopy, we have found new evidence for bar-mediated resonances as the primary drivers of evolution in M94 at the present epoch.   Our observational results include evidence for  \\noindent\\hang(1.)  A 450-pc long nuclear ``mini-bar'' at both optical and near-UV wavelengths.  \\noindent\\hang(2.)  An inner disk with diffuse FUV emission in concentric arcs that do not match the fine-scale structures or reddened structures at visible wavelengths.  Since $H\\alpha$ is observed in absorption here, the UIT/FUV image represents the first view of this non-ionizing but relatively young disk component.  \\noindent\\hang(3.)  UV-Optical colors and spectral indices in the nucleus and inner disk that indicate B and A-type stars in the presence of modest extinction (A$_V$ $\\le$1 mag) along with some LINER activity from the nucleus itself.   \\noindent\\hang(4.)  A 2.2 kpc diameter starbursting ring at the perimeter of the inner disk that is bright at FUV, H$\\alpha$, and radio-continuum wavelengths. The level of starbirth activity in this inner ring rivals the levels observed in starbursting irregular galaxies such as NGC 1569 and NGC 4449.  The inferred star formation rate within the ring and inner disk amounts to 1.5 M$_{\\odot}$ yr$^{-1}$ --- sufficient to build up the stellar mass of the inner disk and bulge in $\\sim$10$^{10}$ yr.  \\noindent\\hang(5.)  No detectable radial offsets between the H$\\alpha$ and FUV rings, thus indicating a 35 km/s speed limit to any outward or inward propagating star formation in the ring, if such a mode is present.  \\noindent\\hang (6.) Two 500-pc size FUV-emitting knots exterior to the ring on diametrically opposite sides of the nucleus. The bisymmetric knots and starburst ring appear to be especially prominent parts of a complex spiral arm structure, as revealed in a spatially-filtered B-band image.  \\noindent\\hang (7.) The starburst ring, bi-symmetric knots, oval disk, and outer pseudo-ring as signposts of resonant dynamics in the disk of M94.  More specifically, the radii of these features match those of various orbital resonances, given a pattern speed of 35 km s$^{-1}$ kpc$^{-1}$ at our adopted distance and inclination. These orbital resonances are most likely driven by some combination of the nuclear mini-bar and oval distortion in the disk.  \\noindent\\hang (8.) A shallow minimum of gravitational stability at the radius of the starburst ring that extends inward into the inner disk. Although too broad to explain the discrete starburst ring, the shallow minimum may help to explain the 10$^7-10^8$-yr old stellar population interior to the ring. \\vskip 12pt  Although we can set a limit on the speed of outward or inward propagating star formation in the ring, we cannot preclude the existence of such a mode.  At a propagation speed of 35 km/s, a wave initiated in the nucleus could traverse the inner disk to the radius of the current starburst ring in only 31 Myr.  Therefore, it is possible that the $10^7-10^8$-yr old stellar population detected in the inner disk is the result of such an outward propagating wave.  The striking difference in emission morphologies at FUV and red wavelengths provides further support for the starburst ring being a transient phenomenon which does not persist at any one radius for very long.  Either these resonant phenomena come and go, as the oval distortions undergo secular evolution, or their operating radii migrate in response to other dynamical influences on the stars and gas (cf. \\markcite{Combes94}Combes 1994; \\markcite{Combes95}Combes et al. 1995; \\markcite{Friedli95}Friedli \\& Benz 1995).  Otherwise, one must invoke strong radial inflows of stars from the starburst ring to populate the inner disk and bulge, a feat requiring unusual circumstances -- e.g. mergers.  The results reported herein may have important implications with regard to observations of the most distant observable galaxies.  At redshifts of 1--5, the 2-kpc diameter starburst ring in M94 would subtend angles of only (0.7$''$ -- 1.0$''$)H$_{\\circ}$/75 in an Einstein-de Sitter Universe (q$_{\\circ}$ = 1/2) and (0.3$''$ -- 0.2$''$)H$_{\\circ}$/75 in an open (Milne) universe (q$_{\\circ}$ = 0) (cf. \\markcite{Narlikar83}Narlikar 1983). The UV-bright nuclear rings evident in NGC 1097, NGC 1317, NGC 1433, NGC 1512, NGC 2997, NGC 4321, and NGC 5248 (\\markcite{maoz95}Maoz et al. 1995; \\markcite{maoz96}Maoz et al. 1996; \\markcite{kuchinsky00}Kuchinsky et al. 2000; \\markcite{marcum00}Marcum et al. 2000) would subtend even smaller angles at the same redshifts.  Moreover, nuclear rings tend to have higher FUV surface brightnesses than their larger counterparts -- the inner ring in M94 being a remarkable exception. Therefore, some of the ``core-halo'' morphologies that are evident at high-redshift in the restframe FUV (cf. \\markcite{Giavalisco97}Giavalisco et al. 1997) may, in fact be marginally-resolved representations of galaxies with starburst rings in their centers.  Gravitationally-lensed galaxies are fortuitously magnified, enabling resolutions of their structure at high S/N.  An important precedent in this regard is the gravitationally-lensed ``Pretzel Galaxy'' which lies behind the galaxy cluster 0024+1654 at an estimated redshift of 1.2 -- 1.8 (\\markcite{Colley96}Colley et al. 1996; \\markcite{Tyson97}Tyson et al. 1997). Detailed reconstructions of the multiply-lensed galaxy show a clear annular morphology on a scale of several kpc.  If M94 and other nearby ringed galaxies can be used as current-epoch analogues, the ``Pretzel Galaxy'' and perhaps other marginally-resolved ``core-halo'' galaxies at high redshift may represent youthful inner disks and bulges growing under the organizing influence of oval or bar asymmetries(\\markcite{Friedli95}Friedli \\& Benz 1995; \\markcite{Waller97}Waller et al. 1997). Conversely, if evidence for starburst rings at high redshift proves to be sparse, then massive inner disks featuring ring-bar dynamics have yet to form in most systems, or starbursting bulges are masking their presence."
    },
    "0011/astro-ph0011323_arXiv.txt": {
        "abstract": "{ The \\object{Coma} cluster of galaxies was observed with XMM-Newton in 12 partially overlapping pointings. We present here the resulting X-ray map in different energy bands and discuss the large scale structure of this cluster. Many point sources were found throughout the observed area, at least 11 of them are coincident with bright galaxies. We also give a hardness ratio map at the so far highest angular resolution obtained for a cluster of galaxies. In this map we found soft regions at the position of bright galaxies, little variation in the central 15 arcmin, but some harder regions north of the line \\object{NGC 4874} -- \\object{NGC 4889}. ",
        "introduction": "The \\object{Coma} cluster of galaxies was one of the performance verification targets of XMM-Newton. The aim was to verify XMM-Newton's ability to map large extended X-ray sources. This is not an easy task because errors can occur in several areas.  The telescope vignetting plays a major role when overlapping the same region of an extended source, measured at different off-axis angles. Because of the vignetting, the surface brightness of a region measured off axis seems to be lower than the same region measured on axis. The reason for that is the increasing loss of mirror collecting area with increasing off-axis angle, which in addition is energy dependent. By de-vignetting the different observations and normalizing the observations to the obser\\-ving times the surface brightnesses of the same source regions are corrected and must show the same number of counts/arcmin$^2$/sec. If done correctly, detector-edges or regions of detector gaps should not be visible as inhomogenities in the surface brightness of the merged observations.  \\begin{table*}[t] \\caption[ ]{Journal of Observations} \\begin{tabular}{llccccc} \\hline \\rule{0mm}{4mm} & Name of     &          &           &  \\multicolumn{3}{c}{pn-Camera observing times (ksec)}\\\\ \\multicolumn{1}{c}{Date} & Observation & RA(2000) & DEC(2000) & planned     & \\multicolumn{1}{c}{performed} & \\multicolumn{1}{c}{effective} \\\\ \\hline \\rule{0mm}{4mm} 2000 May 29 \\    & Coma center & 12 59 46.7 & 27 57 00 & 15.0 & 22.7 & 15.4 \\\\ 2000 June 21/22  & Coma 1      & 12 56 47.7 & 27 24 07 & 25.0 & 29.2 & 24.0 \\\\ 2000 June 11/12  & Coma 2      & 12 57 42.5 & 27 43 38 & 25.0 & 34.8 & 34.7 \\\\ 2000 June 27/28  & Coma 3      & 12 58 32.2 & 27 24 12 & 25.0 & 27.0 & 14.3 \\\\ 2000 June 23     & Coma 4      & 13 00 04.6 & 27 31 24 & 25.0 & 20.0 & (11.4) \\\\ 2000 May 29      & Coma 5      & 12 59 27.5 & 27 46 53 & 20.0 & 25.0 & 10.0 \\\\ 2000 June 12     & Coma 6      & 12 58 50.0 & 27 58 52 & 20.0 & 14.8 &  3.7 \\\\ 2000 June 24     & Coma 7      & 12 57 27.7 & 28 08 41 & 25.0 &  8.1 & (3.4) \\\\ not observed yet & Coma 8      & 13 01 25.6 & 27 43 53 & 26.0 &   -- &   -- \\\\ 2000 June 12     & Coma 9      & 13 00 32.7 & 27 56 59 & 20.0 & 26.3 & 21.3 \\\\ 2000 June 22     & Coma 10     & 12 59 38.4 & 28 07 40 & 20.0 & 26.3 & 16.6 \\\\ 2000 June 24     & Coma 11     & 12 58 36.5 & 28 23 56 & 25.0 & 32.1 & 15.9 \\\\ not observed yet & Coma 12     & 13 01 50.2 & 28 09 28 & 25.0 &   -- &   -- \\\\ not observed yet & Coma 13     & 13 00 36.5 & 28 25 15 & 25.0 &   -- &   -- \\\\ 2000 June 22     & Coma bkgd   & 13 01 37.0 & 27 19 52 & 30.0 & 22.7 & 20.1 \\\\ \\hline \\end{tabular} \\end{table*}     The backgrounds in the image are problematic.  The part of the background that is produced by particles inside the detector is not distinguishable from X-rays, is not vignetted, and is in general distributed homogeneously over the detector face (except for the lower end of the spectral range and around the detector intrinsic emission lines around 8 keV, see Briel et al. \\cite{ub00}). Hence this component must be subtracted before de-vignetting is applied. The extragalactic X-ray background is vignetted in the same manner as diffuse emission from an extended source and may therefore be de-vignetted in the same manner. But there is a third component in the background, the low energy protons, that are again indistinguishable from X-rays. They come through the mirror system and show a different vignetting compared to X-rays (see Str\\\"uder et al. \\cite{ls01}). In addition, their spectrum and their flux is time variable. Because of the different vignetting, it is essential to subtract a ``proton image\" from the measured image before de-vignetting is applied, especially when the surface brightness of the proton induced background is in the same order of magnitude as the surface brightness of the extended source.  \\object{Coma} was chosen because it is a well studied source at all wavelength ranges. With a diameter of the main cluster, which is several times larger than the FOV of the EPIC cameras of about 26 arcmin, it is necessary to mosaic the cluster with XMM-Newton if the large scale structure and dynamics of the cluster are to be investigated. Although \\object{Coma} was long believed to be the archetype of a relaxed cluster of galaxies, the X-ray image taken during the Rosat All Sky Survey revealed the complex large scale structure, especially the subgroup around \\object{NGC 4839} probably falling into the main cluster (Briel et al. \\cite{ub92}). Subsequent pointed observations with Rosat showed that \\object{Coma} has irregular structure on different scales and was most probably formed by hierarchical clustering of several subunits, which are still visible in their X-ray emission (White et al. \\cite{sw93}). Observations with ASCA, and also the pointed Rosat observation, showed evidence of temperature structure in \\object{Coma} (Honda et al. \\cite{hh96}; Briel \\& Henry \\cite{ub98}; Donnelly et al. \\cite{rd99}; Watanabe et al. \\cite{mw99}). The interpretation of the temperature structure favors the ongoing merging  of the subgroup around \\object{NGC 4839} with the main cluster, a scenario which is still debated.  In this letter we will present the result of merging 12 partially overlapping observations of the \\object{Coma} cluster with the EPIC pn-CCD-detector on board XMM-Newton. In Section 2 we will describe the observations and present the data reduction and the method of merging the individual observations. In Section 3 we will discuss some of the many point sources in the image, the large scale structure of the X-ray surface brightness distribution, and the hardness ratio map of the inner cluster region. ",
        "conclusions": "The \\object{Coma} cluster is the nearest very rich cluster of galaxies. It is probably the best studied cluster at all wavelengths and is one of the brightest extragalactic X-ray sources in the sky. Until now, the deepest X-ray observation of \\object{Coma} is that of White et al. (\\cite{sw93}) consisting of 4 exposures of about 19 ksec each by the ROSAT PSPC.  The EPIC pn observations reported here are of comparable length, thus are $\\sim$8 times deeper, while simultaneously extending over $\\sim$3 times wider energy range and with $\\sim$4 times better angular resolution.  Perhaps not surprisingly given the better spatial re\\-solution and deeper observation, the most noticable difference between our XMM images in Figures 1 and 2 compared to Figure 2 of White et al. (\\cite{sw93}) is the myriad of newly detected point-like sources. Dow \\& White (\\cite{kd95}) discuss the $\\sim$25 point sources in the ROSAT observations of \\object{Coma}. There are at least 75 point sources in the smaller solid angle we have analyzed. We have made a first attempt to identify these sources by overlaying our X-ray map on an optical image of the region and then consulting NED when there appears to be a coincidence. Many bright galaxies seem to be X-ray sources including: \\object{NGC 4827}, \\object{4839}, \\object{4858}, \\object{4860}, \\object{4874}, \\object{4889}, \\object{4911}, and \\object{4921}, and \\object{IC 4040} and \\object{4051}. However some faint galaxies, with magnitudes from 18 to 20, could also be identifications based on positional coincidence. At least three quasars, \\object{QSO 1256+281}, \\object{QSO 1258+280},  and \\object{QSO 1259+281}, are detected. However, the bulk of the sources are unidentified in our initial analysis.  The \\object{NGC 4911} and \\object{NGC 4839} groups contain concentrations of point sources at our sensitivity. Dow \\& White (\\cite{kd95}) and Neumann et al. (\\cite{dn01}) discuss the particularly interesting morphology of the \\object{NGC 4839} group.  The large-scale diffuse emission in the soft band shown in Figure 1 is very similar to the ROSAT observations of White et al. (\\cite{sw93}) in the same band. The X-ray emission of the \\object{Coma} cluster is lumpy near its center but becomes smoother at larger radii. An additional lump, apparent in the ROSAT observation but not discussed by White et al. (\\cite{sw93}), is directly NE of \\object{NGC 4839}. It may be yet another group on its way into the main cluster to the ENE. This lump may be preceeding the \\object{NGC 4839} group from the direction of \\object{A1367} since calculations of cluster growth indicate that the groups that merge to form clusters flow along filaments connecting the clusters. The large-scale diffuse hard emission is very similar to the soft, as Figure 2 shows.  Because the calibration of the spectral response of the extended full frame mode of the pn-CCD is not yet available, we do not yet want to perform spatially resolved spectral fits. We can, however, make a hardness ratio map using wide bands, which we show in Figure 3. While these maps sacrifice spectral details, such as temperatures and abundances, they are the highest angular ($\\sim$40 arcsec) and spatial ($\\sim$26 kpc for H$_{\\mbox{o}}$ = 50 km s$^{-1}$ Mpc$^{-1}$) resolution spectral maps ever obtained for a cluster of galaxies.  Using the conversion from hardness ratio to temperature, given in Fig. 3, we have checked our resulting mean temperature in some regions which are known from other measurements. For the over-all temperature within a radius of 10 arcmin, centered on \\object{NGC 4874}, we determined a temperature of 8.8 keV. The GINGA value is 8.21 $\\pm$ 0.09 keV (Hughes et al. \\cite{jh93}), the ASCA value is 9. $\\pm$ 0.6 keV (Donnely et al \\cite{rd99}) and the XMM-MOS value is 8.25 $\\pm$ 0.10 keV (Arnaud et al \\cite{ma01}). We agree to better than 10\\% with these previously determined values, which would be expected given our understanding of the uncertainties of the calibration and of the background model. In a further check, we determined the temperature around \\object{NGC 4874} within a radius of 0.5 arcmin and within a ring from 1 to 5 arcmin. Our result is 7 and 9.2 keV, the MOS result (Arnaud et al \\cite{ma01}, Fig. 8) is 6.6 and 8.6 keV respectively, again in agreement within 10\\%. We found the same kind of agreement on other places, taking into account the large area of 3.5 $\\times$ 3.5 arcmin$^2$, used for the MOS data by Arnaud et al. (\\cite{ma01}).   The large-scale description of the data in Figure 3 agrees with Arnaud et al. (\\cite{ma01}): there is cool gas to the south-east and hotter gas to the south-west and west. It agrees to some extent with Donnely et al. (\\cite{rd99}), Watanabe et al. (\\cite{mw99}) and Briel \\& Henry (\\cite{ub98}) investigations. In particular the south-east cool region  was already identified by Donnely et al. (\\cite{rd99}) and a hot region in the south-west/west direction, but somewhat further away, was found by Watanabe et al. (\\cite{mw99}).  There are spectral variations on the smallest scales that we can measure. Softer emission is associated with many compact sources some of which are associated with galaxies as discussed previously. Although the smoothing of the original surface brightness distribution smears out the point sources to some extend, a significant drop in the radial temperature  profile within 1 arcmin of \\object{NGC 4874} was also seen by Arnaud et al. (\\cite{ma01}), Fig. 8. The Rosat temperature map of Briel \\& Henry (\\cite{ub98}) also shows cooler regions around the two central galaxies (although other cool regions in the center of their map are not confirmed).  If these soft regions, associated with galaxies, are extended, we can postulate three possible explanations. First, they are local potential wells centered on cluster galaxies as is was discussed by Vikhlinin et al. (\\cite{av94}). The shallower potential wells could only trap lower temperature gas. Second, they are the interstellar emission that the cluster galaxy has been able to retain. Third, integrated emission from low mass X-ray binaries which Irwin \\& Sarazin (\\cite{ji98}) have discussed as the source of soft emission from elliptical galaxies.  Perhaps the most interesting small scale spectral structures are the isolated hard spots that do not appear to be associated with any obvious point-like X-ray emission.  For the most part they are not associated with an optical object in NED. The five hard spots north of the \\object{NGC 4889} -- \\object{NGC 4874} line may be the cause of the hot region in the same general location detected by Donnelly et al. (\\cite{rd99})  in their ASCA observations. It seems that these hot spots are washed out in the temperature map of Arnaud et al. (\\cite{ma01}) because of the larger (3.5$\\times$3.5 armin$^2$) integration boxes. They might have been seen however in the even  larger integration region of ASCA, if they are associated with variable sources like QSO's. But we can not entirely eliminate an instrumental cause  for some of them because some are aligned parallel to the CCD orientation and on the boarder of the chips."
    },
    "0011/astro-ph0011115_arXiv.txt": {
        "abstract": "A search for star forming objects belonging to tidal tails has been carried out in a sample of deep H$\\alpha$ images of 16 compact groups of galaxies. A total of 36 objects with H$\\alpha$ luminosity larger than 10$^{38}~erg~s^{-1}$ have been detected in five groups. The fraction of the total H$\\alpha$ luminosity of their respective parent galaxies shown by the tidal objects is always below 5\\% except for the tidal features of HCG~95, whose H$\\alpha$ luminosity amounts to 65\\% of the total luminosity. Out of this 36 objects, 9 star forming tidal dwarf galaxy candidates have been finally identified on the basis of their projected distances to the nuclei of the parent galaxies and their total H$\\alpha$ luminosities. Overall, the observed properties of the candidates resemble those previously reported for the so-called tidal dwarf galaxies. ",
        "introduction": "Interactions between disk galaxies are known to produce strong morphological disturbances and, if the geometry of the encounter is favorable, to enhance the star formation rates of the disk galaxies. The pioneering work by Tomre \\& Tomre (1972) and a later study by Barnes (1988) supported the original idea by Zwicky (1956) that dwarf galaxies could originate within the tidal tails resulting from a strong interaction between two disk galaxies. Since then, several examples of dwarf galaxies with tidal origin have already been detected in interacting systems in an advanced stage of merging and also around galaxies experiencing less dramatic interactions (Schweizer 1978, Mirabel et al. 1991,1992, Hibbard \\& van Gorkom 1996, Duc \\& Mirabel 1997,1998).  The compact group environment has been suggested to be an ideal place for interactions between galaxies given their large spatial densities and their low velocity dispersions. Mendes de Oliveira \\& Hickson (1994) reported morphological signs of interaction in many galaxies of compact groups from Hickson's Catalogue (1982). Other recent studies on the photometric properties of the galaxies in Hickson Compact Groups at different wavelengths resulted contradictory when compared to samples of field galaxies: while CO, FIR, H$\\alpha$ and optical fluxes seem to be similar to those measured for field galaxies (Moles et al. 1994, Verdes-Montenegro et al. 1998, Iglesias-P\\'{a}ramo \\& V\\'{\\i}lchez 1999), radio continuum and 21cm emission are on average lower than for samples of isolated galaxies (Williams \\& Rood 1987, Menon 1995, Huchtmeier 1997).  Hunsberger et al. (1998) carried out a search for dwarf galaxy candidates in the close environment of a sample of Hickson Compact Groups, by automated detection of faint objects on deep $R$ band frames. After decontamination for the presence of background galaxies, they found that the luminosity function in compact groups showed an enhancement at the faint luminosity end compared to a normal Schechter function. They concluded that the initial dwarf galaxy population in compact groups is replenished by ``subsequent generations'' formed in the tidal debris of giant galaxy interactions.  In this work we present the results of a search for tidal star forming dwarfs in our H$\\alpha$ images of a sample of nearby compact groups (see V\\'{\\i}lchez \\& Iglesias-P\\'{a}ramo 1998, Iglesias-P\\'{a}ramo \\& V\\'{\\i}lchez 1999). Five of the groups were found to show optically prominent tidal tails with star forming objects within them. We emphasize the fact that all the detected objects are H$\\alpha$ emitters, therefore implying that they are physically associated with the galaxies of the group. The main properties of the star forming objects will be presented in \\S\\ref{sample}, as well as a description of their close environment. \\S\\ref{discussion} contains a discussion on their likely evolutionary scenario. ",
        "conclusions": "}  As can be seen from Table~1, the H$\\alpha$ equivalent widths and luminosities of the tidal objects suggest that most of them host very young star formation bursts and are above the luminosity limit for giant H{\\sc ii} regions. The underlying continuum luminosity varies greatly from object to object as shown in Figures~1 to 5. This is the reason for the large range covered by their H$\\alpha$ equivalent width values, rather than the result of different evolutionary stages of the corresponding star formation bursts of the tidal objects. The fraction of the total H$\\alpha$ flux corresponding to tidal objects is typically under 5\\% (see Table~1), except in the remarkable case of the tidal features in HCG~95c, which amount to 65\\% of the total luminosity of the galaxy. We attribute this extraordinarily high fraction of luminosity within the tidal tail to the fact that HCG~95c belongs to a triple system showing a rather complicated interaction sequence (Iglesias-P\\'{a}ramo \\& V\\'{\\i}lchez 1998).  The existence of these tidal objects is linked to the possibility that they could subsequently evolve as independent dwarf galaxies, like those observed at the tips of the tidal tails of the Antennae and the Superantennae (Mirabel et al. 1991,1992). Looking to the optical images, only rough morphological similarities between the tails of the Antennae and Superantennae and the tails shown by our galaxies can be found. The Antennae and Superantennae fall in the most dramatic case of interaction of a pair of galaxies, whereas our five groups span a wide range of interaction strengths. At this point, two critical questions arise concerning the possibility that tidal objects can evolve independently as dwarf galaxies: (1) are they distant enough from the parent galaxy so as to be able to escape from the potential well?, and (2) are they massive enough to be self gravitating systems and to be considered dwarf galaxies?  Concerning the first question, some authors have claimed (Schweizer 1978; Hibbard \\& van Gorkom 1996) that most of the material ejected during an interaction is accumulated in the tidal tails. N-body simulations suggest that a substantial fraction of this material will slowly fall into the main body of the remnant of the interaction, and only the outermost 20\\% will probably gain enough kinetic energy to be able to evolve independently for a long time (see Hibbard \\& Mihos 1995). Given that no information is available on the velocity fields of the interacting pairs of our sample, and that they show a great variety of interaction patterns, we will adopt a quantitative criterion based on the projected distances of the tidal objects from the nuclei of their parent galaxies in order to select the most likely candidates to escape from the potential well. Ferguson et al. (1998) carried out a deep search for H{\\sc ii} regions well outside the optical limits of a sample of disk galaxies, and found no H{\\sc ii} regions located at a projected galactocentric distance larger than $2 \\times R_{25}$. This result was explained as due to the fact that the gas surface density is below the critical density for star formation at such large radii from the centers of the galaxies. Therefore, we can consider that those star forming regions located that far from the nuclei of their parent galaxies must have been originated by an external mechanism which has removed gas from the inner regions of the parent galaxies. Star forming regions far away from their parent galaxies may become kinematically decoupled from the stellar disks; thus they are very likely escaping the potential well of the parent galaxy and could become good tidal dwarf galaxy candidates. According to Table~\\ref{prop}, {\\em only 9 tidal objects fulfill this criterion} ($R > 2 \\times R_{25}$). Obviously, this is a conservative criterion adopted given the lack of information on the velocity field of the tidal tails. Therefore, we should bear in mind that because of possible projection effects, the number of tidal dwarfs candidates selected is a lower limit since other objects not satisfying this criterion cannot be ruled out to be tidal dwarf galaxies.  Concerning the second question posed above, Elmegreen et al. (1993) proposed that clouds as massive as $10^{8}~M_{\\odot}$ can form during interactions of disk galaxies. Under some conditions, these clouds could become dwarf galaxies. In particular, when the mass of the perturber is larger than approximately 1.4 times the mass  of the parent galaxy, -- as appears to be the case for the five groups of our sample presenting H$\\alpha$ emitting tidal tails\\footnote{estimated from a comparison of the magnitudes of the parent galaxies with those of the remaining galaxies in the groups} -- these models predict that the outermost clouds are very likely to escape the potential well. Such $10^{8}~M_{\\odot}$ clouds would produce $10^{6}~M_{\\odot}$ starbursts, assuming a value for the efficiency of star formation of 1\\% (see Kennicutt 1998 and references therein). From the observational side, the H$\\alpha$ luminosities reported for the tidal dwarf galaxies of the Antennae (Mirabel et al. 1991) and NGC~5291 (Duc \\& Mirabel 1998), are comparable to the ones derived in the present paper; the H{\\sc i} masses measured for those dwarfs appear to be typically above $10^{8}~M_{\\odot}$. In the same direction, from a large compilation of dwarf galaxies Hunter \\& Gallagher (1986) reported H{\\sc i} masses larger than $10^{8}~M_{\\odot}$  for objects with $L_{\\mbox{\\scriptsize{H}}\\alpha} \\gtrsim 10^{39}~erg~s^{-1}$. Such bright objects are known to be gravitationally bound, according to the well established $L($H$\\alpha$) {\\em versus} $\\sigma$ relationship for H{\\sc ii} galaxies (Terlevich \\& Melnick 1981). The 9 objects selected above show H$\\alpha$ luminosities larger than $10^{39}~erg~s^{-1}$, which on observational grounds seems to be a reasonable lower limit to ensure self gravitation for H{\\sc ii} complexes.  The two restrictions imposed above, though conservative, appear reasonable since H{\\sc ii} regions brighter than $10^{39}~erg~s^{-1}$ are found to be located at galactocentric distances lower than about $1.2 \\times R_{25}$ (as derived from the H{\\sc ii} region surveys by Rozas et al. 1996,2000). A naive estimation of the dynamical time needed by our dwarf galaxy candidates to travel a distance equivalent to $R_{25}$ as a consequence of the interaction would be of the order of the fading time for H{\\sc ii} regions. Thus, we are confident that the 9 selected candidates must have been generated during the interaction and do not belong to a previous generation of disk H{\\sc ii} regions.  From the arguments presented above, it appears clear that the 9 star forming objects finally selected represent good candidates to satisfy the {\\em escape} and {\\em self gravitation} conditions, and thus they can be included in the group of the so-called tidal dwarf galaxies. Further information on total masses and relative velocity with respect to their parent galaxies is required to confirm that they satisfy both conditions.  It seems striking that the ratio between the number of candidate dwarf galaxies to the total number of tidal objects detected, 25\\%, appears consistent with the fraction of the ejected material which is expected to escape according to theoretical models (Hibbard \\& van Gorkom 1995). Since all these candidate tidal galaxies were selected at  projected distances larger than $2\\times R_{25}$, we suggest that this last criterion might be a useful indicator of an {\\em effective} escape radius for these systems.  Summarizing, as a result of a search performed on a sample of net H$\\alpha$ images of 16 compact groups of galaxies, we have identified 9 star forming tidal objects proposed to be considered dwarf galaxy candidates given their large projected distances from their parent galaxies and their large H$\\alpha$ luminosities. These interesting objects deserve further detailed study, including dynamical and evolutionary aspects which are beyond the scope of the present study."
    },
    "0011/astro-ph0011265_arXiv.txt": {
        "abstract": " ",
        "introduction": "The concept of biasing between different classes of extragalactic objects and the background matter distribution was put forward by Kaiser (1984) and Bardeen et al. (1986) in order to explain the higher amplitude of the 2-point correlation function of clusters of galaxies with respect to that of galaxies themselves.  In this framework biasing is assumed to be statistical in nature; galaxies and clusters are identified as high peaks of an underlying initially Gaussian random density field. Biasing of galaxies with respect to the dark matter distribution was also found to be an essential ingredient of CDM models of galaxy formation in order to reproduce the observed galaxy distribution (cf. Davies et al. 1985; Benson et al. 2000).  The classical approach to study the redshift evolution of bias utilises the ratio of the correlation functions of objects and dark matter, which are assumed to be related via the square of a scale independent bias factor. However, in this study we will use the definition by which the extragalactic mass tracer (galaxies, halos, clusters) fluctuation field, $\\delta_{tr}$, is related to that of the underlying mass, $\\delta_{\\rm m}$, by \\be\\label{eq:1} \\delta_{\\rm tr} = b \\delta_{\\rm m} \\;\\;, \\ee where $b$ is the linear bias factor. Note that the former definition results from the latter but the opposite is not necessarily true. The bias factor may have many dependencies; even assuming that it is scale independent, it necessarily depends on the type of the mass tracer as well as on the epoch $z$, since the fluctuations evolve with time as gravity draws together galaxies and mass. It is evident, therefore, that the bias factor should also depend on the different cosmological models and dark matter content of the Universe (for a recent overview see Klypin 2000).  More realistic biasing schemes have been proposed in the literature. Coles (1993) introduced the idea of biased galaxy formation in which galaxies form with a probability given by an arbitrary function of the local mass density. Mann, Peacock \\& Heavens (1998) investigated the properties of different bias models of galaxy distributions that results from local transformations of the present-day density field. The deterministic and linear nature of eq.(\\ref{eq:1}) has been challenged (cf. Dekel \\& Lahav 1999; Tegmark \\& Bromley 1999) and indeed some non-linearity of the biasing relation is necessary to reconcile high biasing with deep voids. Despite the above, the linear biasing assumption is still a useful first order approximation which, due to its simplicity, it is used in most studies of large scale (linear) dynamics (cf. Strauss \\& Willick 1995 and references therein; Branchini et al. 1999; Schmoldt et al. 1999; Plionis et al. 2000).  Different studies have indeed shown that the bias factor is a monotonically decreasing function of redshift. An important advancement in the analytical treatment of the bias evolution was the work of Mo \\& White (1996) in which they used the Press-Schechter (1974) formalism and found that in an Einstein-de Sitter universe the linear bias factor evolves strongly with redshift. Using a similar formalism, Matarrese et al. (1997) extended the Mo \\& White results to include the effects of different mass scales (see also Catelan et al 1998).  Steidel et al. (1998) confirmed that the Lyman-break galaxies are very strongly biased tracers of mass and they found that $b(z=3.4) \\magcir 6, 4, 2$, for SCDM, $\\Lambda$CDM $(\\Omega=0.3)$ and OCDM $(\\Omega=0.2)$, respectively (see also Giavalisko et al 1998). A similar value for the $\\Lambda$CDM model was obtained by Cen \\& Ostriker (2000) using high resolution Nbody/hydro simulations in which they treated DM, gas as well as star formation. The use of high resolution N-body simulations (cf. Klypin et al. 1996; 1999, Cole et al. 1997 and references therein) have shown that anti-biasing ($b< 1$) should exist at scales $r \\sim 3 - 8 h^{-1}$ Mpc, for the open and flat low-$\\Omega$ models, in contrast with $\\Omega=1$ models, where $b>1$. Colin et al (1999), using high-resolution N-body simulations of SCDM, $\\Lambda$CDM, OCDM and $\\tau$CDM models, which avoid the so called ``overmerging\" problem, found that indeed biasing evolves rapidly with redshift, while Kauffmann et al. (1999) combining semi-analytic models of galaxy formation and N-body simulations has also studied the evolution of clustering in different cosmologies.  In this paper we will not indulge in such aspects of the problem but rather, working within the paradigm of linear and scale-independent bias, we will derive the functional form of its redshift evolution in the matter dominated epoch and in all cosmological models. The Einstein de-Sitter case has been studied in the past (cf. Nusser \\& Davis 1994; Fry 1996; Bagla 1998) using the continuity equation, which is a first order differential equation, to derive a solution, $\\propto (1+z)$, valid only for low $z$'s. Our approach is to use the perturbation evolution equation which combines the continuity, the Euler and the Poisson equations and which is a second order differential equation. We should therefore expect to find a further component to the known solution.  The paper is organised as follows: in section 2 we discuss the basic models for the linear bias evolution, in section 3 we derive the basic differential equation describing the evolution of the linear bias factor, while in section 4 we present its analytical solution for the different cosmological models and a comparison with previous models and N-body simulation results. Finally, in section 5 we summarise our main results. ",
        "conclusions": "We have introduced analytical arguments and approximations based on linear perturbation theory and a linear, scale-independent bias between a mass tracer and its underlying matter fluctuation field in order to investigate the cosmological evolution of such a bias. We derive a second order differential equation, the solution of which provides the functional form of the of bias evolution in any Universe. For the case of an Einstein-de Sitter Universe, we find an exact solution which is a linear combination of the known solution $\\propto (1+z)$ (cf. Bagla 1998 and references therein), derived from the continuity equation, and a second term $\\propto (1+z)^{3/2}$ which dominates. This solution once parametrised at two different epochs, compares extremely well with the more sophisticated halo models (cf. Mo \\& White 1996) and with N-body simulations.  For the two low-density cosmological models we find exact solutions, albeit only in the high-redshift approximation (where the growing mode of perturbations can be approximated by the Einstein-de Sitter solution). We also derive analytical solutions for two limiting low-density Universes (ie., $\\Omega=0$, $\\Omega_{\\Lambda}=0$ and $\\Omega=0$, $\\Omega_{\\Lambda}=1$)."
    },
    "0011/astro-ph0011053_arXiv.txt": {
        "abstract": "OO Aql is a rare W UMa-type eclipsing binary in which the two solar-type stars may have only recently evolved into contact. The binary has an unusually high mass ratio (0.84), and a relatively long orbital period (0$\\fd$506) for it spectral type (mid-G).  Twelve ultraviolet spectra of OO Aql were obtained in 1988 with the {\\it IUE} satellite, including a series of consecutive observations that cover nearly a complete orbital cycle.  Chromospheric activity is studied by means of the Mg~{\\sc ii} h$+$k emission at 2800 {\\AA}. The Mg~{\\sc ii} emission is found to vary, even when the emission is normalized to the adjacent continuum flux.  This variation may be correlated with orbital phase in the 1988 observations.  It also appears that the normalized Mg~{\\sc ii} emission varies with time, as seen in spectra obtained at two different epochs in 1988 and when compared with two spectra obtained several years earlier.  The level of chromospheric activity in OO Aql is less than that of other W UMa-type binaries of similar colors, but this is attributed to its early stage of contact binary evolution.  Ultraviolet light curves were composed from measurements of the ultraviolet continuum in the spectra.  These were analyzed along with visible light curves of OO Aql to determine the system parameters.  The large wavelength range in the light curves enabled a well-constrained fit to a cool spot in the system. ",
        "introduction": "The eclipsing binary star system OO Aquilae (HD 187183, $V_{\\rm max}=9.2$, $B-V=+0.76$, mid-G) is an unusual W UMa-type binary because it possesses a mass ratio near unity (0.84; Hrivnak 1989).  This suggests that the components have only recently (on astronomical timescales) come into contact, since once in contact the stars are expected to evolve toward low mass ratio (Webbink 1976; Vilhu 1982). The orbital period of 0$\\fd$506 is long compared to other G spectral types stars with W UMa-type light curves.  This results in an unusually large angular momentum for a cool W UMa-type system, again suggesting that it has only recently evolved into contact (Mochnacki 1981).  Thus OO Aql represents a rare, transient phase in the evolution of contact binary star systems.  On the basis of his modern radial velocity study using the cross-correlation technique, Hrivnak (1989) carried out a consistent analysis of both the light and velocity curves of OO Aql.  He determined the absolute parameters of the two stars with high precision (M$_1$=1.04$\\pm$0.02 M$_{\\sun}$, R$_1$=1.39$\\pm$0.02 R$_{\\sun}$, M$_2$=0.88$\\pm$0.02 M$_{\\sun}$, R$_2$=1.29$\\pm$0.02 R$_{\\sun}$).  By comparing its properties with stellar models, Hrivnak presented evidence that the stars have an age of $\\sim$8 Gyr.  Thus OO Aql consists of a pair of mid-G stars in contact, which appear to be somewhat more evolved than the Sun.  It falls into the subclass of A-type W UMa-type binaries, in which the more massive component is eclipsed at primary minimum; this is unusual among the cooler (G and K spectral type) W UMa-type binaries.  W UMa-type binaries of G and K spectral type are known to display a high level of chromospheric activity, as shown in the study of Rucinski (1985). Most contact binaries of G$-$K spectral types have periods of 0$\\fd$25 to 0$\\fd$35 and mass ratios of  0.3 to 0.5.  Thus OO Aql is clearly unusual.  Another contact binary with a large mass ratio (0.94) is VZ Psc (P=0$\\fd$26, K2$-$5). This system displays a high level of chromospheric activity, with strong and variable Mg~{\\sc ii} h+k and Ca~{\\sc ii} H+K emission (Hrivnak, Guinan, \\& Lu 1995).  This suggests that OO Aql might also have variable chromospheric activity.   In OO Aql, Ca~{\\sc ii} emission is not evident in medium-resolution spectra, although there may be infilling of the broad absorption profiles (Hrivnak 1989).  However, this apparent absence is most probably the result of a contrast effect, since OO Aql is hotter and has a much higher continuum in this spectral region (3900$-$4000 {\\AA}) than does VZ Psc.  Since the continuum is lower in the ultraviolet, the Mg~{\\sc ii} h+k emission at 2800 {\\AA} provides a better opportunity to measure the chromospheric activity in OO Aql.  A single observation of OO Aql made with the low-dispersion LWR camera of the {\\it International Ultraviolet Explorer} ({\\it IUE}) satellite by Rucinski (1985) in 1984 indicated an unusually low Mg~{\\sc ii} emission level when compared with other W UMa-type systems.  No x-ray detections, which would measure the coronal activity, have been reported for OO Aql, not even in the ROSAT Faint Source Catalogue (Voges et al. 2000).  However, it has been measured as a variable radio source.  It was one of only a few systems detected in a 3.6-cm radio continuum survey of W UMa-type binaries, but was not detected at similar sensitivity when observed two years later (Rucinski 1995).  With this background of well-determined but unusual properties (compared with other W UMa-type binary systems), we were motivated to carry out an ultraviolet study of OO Aql with the {\\it IUE} satellite during a complete orbital cycle.  Our goal was to investigate the level of its Mg~{\\sc ii} emission and to search for phase-related and time-related variations.  As a by-product of these spectra, we also obtained ultraviolet light curves.  In this paper we present the results of this study. ",
        "conclusions": "Ultraviolet observations were obtained for the short-period contact binary OO Aql on two dates in 1988, including nine consecutive observations that covered most of one orbital cycle.  Mg~{\\sc ii} h+k emission is seen in each of the spectra.  These were measured for emission strength, and also measurements were made of the continuum at several wavelength intervals and used to compose ultraviolet light curves for the binary. An analysis of the Mg~{\\sc ii} emission indicates that real variations in the strength of the feature exist for OO Aql, even when normalized to the surrounding ultraviolet continuum.  These may be correlated with the orbital phase, or they may instead indicate overall changes in the chromospheric activity level with time.  Previous studies of phase-related activity in W UMa-type binaries have yielded similar results.  Eaton (1986) found only a slight phase dependence in the normalize Mg~{\\sc ii} emission for SW Lac (P=0$\\fd$32, K0), another contact binary with a relatively large mass ratio (0.73, Hrivnak 1992).   The study of VZ Psc (Hrivnak et al. 1995) revealed variations in the normalized Mg~{\\sc ii} emission that were not phase dependent, but that seemed to vary in time. VW Cep (P=0$\\fd$27, G5V+G8V) shows primarily time-dependent variations in chromospheric Mg~{\\sc ii} emission and also in the transition region emission flux (Guinan \\& Gim\\'enez 1993).  Thus variations in the normalized Mg~{\\sc ii} emission are common in W UMa-type binaries, but these are typically not phase dependent, and instead are probably due to variations in chromospheric activity arising from possible flaring.  General trends have been found for W UMa-type and other short-period binaries in studies of chromospheric activity versus color or orbital period.  A summary for W UMa-type binaries has been published by Rucinski (1985), which included his observation of OO Aql.  While these more recent observations show OO Aql to have a somewhat higher level of chromospheric activity than at the time of his 1984 observation, the basic interpretation is the same; OO Aql has a much lower level of relative activity (by a factor of 2 relative to f(bol)) than do other contact binaries with similar colors.  However, this is presumably due to the recent evolution of its components into contact. While in contact, OO Aql is expected to transfer mass from the less-massive secondary to the more-massive primary component (Webbink 1976), reducing the mass ratio and causing the system to evolve to evolve to an F spectral type. The level of chromospheric activity in OO Aql is in fact comparable to that in W UMa-type systems with the colors of F stars.  Ultraviolet light curves are presented here for the first time for OO Aql. They show a similar shape and amplitude as the visible light curves.  Since the two stars have similar temperatures, there is expected to be little change in the light curves with wavelength. Synthetic light curve solutions were carried out for $B$ and $V$ light curves for two different epochs, and similar orbital and physical parameters were obtained.  These were used to model the new ultraviolet light curves, and good fits were obtained. The combination of ultraviolet and visible light curves led to a well-determined fit for a cool spot on the primary component.  It appears that the light curves at the different epochs can each be fit by the same orbital and physical parameters, with the presence of one major cool spot on the primary component that varies in longitude and size with time. Jeong et al. (1994) carried out a somewhat similar ultraviolet and visible light curve analysis for SW Lac, also including a single cool spot.  The high mass ratio and relatively long orbital period lead to a relatively large orbital angular momentum for OO Aql, which is  larger than that of most contact binaries but less than that of the short-period detached binaries with G-type components, such as ER Vul and UV Leo.  It is likely that the detached systems such as ER Vul (P=0$\\fd$70, M$_1$=0.96 M$_{\\sun}$, M$_2$=0.89 M$_{\\sun}$, Budding \\& Zeilik 1987) and UV Leo (P=0$\\fd$60, M$_1$=0.99 M$_{\\sun}$, M$_2$=0.92 M$_{\\sun}$, Popper 1980) will evolve into contact systems like OO Aql (P=0$\\fd$51, M$_1$=1.05 M$_{\\sun}$, M$_2$=0.88 M$_{\\sun}$) through angular momentum loss by magnetic breaking (Mochnacki 1981; Vilhu 1982).  However, ER Vul has a much higher level of variable starspot and chromospheric activity than that seen in OO Aql, even though its rotational period is longer and its spectral type earlier than those of OO Aql. If OO Aql is typical of a system of solar mass stars that have recently evolved into contact, then the previous high level of activity must be significantly reduced in the contact phase.  Subsequent evolution while in contact is expected to decrease the mass ratio, and if this is so, then a binary like OO Aql will gradually evolve from G to F spectral type as the mass is transferred from the secondary to the primary component.  As shown by Rucinski (1985), the F spectral type systems show a lower activity, consistent with what is now seen in OO Aql.  Thus this study of OO Aql has presented an opportunity to learn more about the properties of a close binary in a rather short-lived stage of its evolution, shortly after the two stars have evolved into contact.  If OO Aql is typical of this stage, then it appears that the level of chromospheric activity must be reduced rather quickly upon entering the contact configuration, with the contact binary continuing to maintain this lower level with some smaller variations over time."
    },
    "0011/astro-ph0011029_arXiv.txt": {
        "abstract": "The declining lightcurve of the optical afterglow of gamma-ray burst GRB000301C showed rapid variability with one particularly bright feature at about $t$$-t_0$=3.8 days.  This event was interpreted as gravitational microlensing by Garnavich, Loeb \\& Stanek (2000) and subsequently used to derive constraints on the structure of the GRB optical afterglow. In this paper, we use these structural parameters to calculate the probability of such a microlensing event in a realistic scenario, where all compact objects in the universe are associated with observable galaxies.  For GRB000301C at a redshift of $z$=2.04, the {\\sl a posteriori} probability for a microlensing event with an amplitude of $\\Delta m$$\\ge$0.95\\,mag (as observed) is 0.7\\% (2.7\\%) for the most plausible scenario of a flat $\\Lambda$-dominated FRW universe with $\\Omega_{\\rm m}$=0.3 and a fraction $f_*$=0.2 (1.0) of dark-matter in the form of compact objects. If we lower the magnification threshold to $\\Delta m$$\\ge$$0.10$\\,mag, the probabilities for microlensing events of GRB afterglows increase to 17\\% (57\\%). We emphasise that this low probability for a microlensing signature of almost a magnitude does {\\sl not} exclude that the observed event in the afterglow lightcurve of GRB000301C was caused by microlensing, especially in light of the fact that a galaxy was found within 2~arcsec from the GRB. In that case, however, a more robust upper limit on the {\\sl a posteriori} probability of $\\approx$5\\% is found. It does show, however, that it will not be easy to create a large sample of strong GRB afterglow microlensing events for statistical studies of their physical conditions on micro-arcsec scales. ",
        "introduction": "Gravitational microlensing offers a way to study the structure of high redshift sources on micro-arcsecond scales, besides being able to constrain the mass fraction and mass function of compact objects in the universe. In strong gravitational lens systems (i.e. systems with multiple images of a single background source), the microlensing optical depth is of order unity (e.g. Chang \\& Refsdal 1979; Gott 1981; Young 1981). Precisely because of this high optical depth, these systems are perfect for resolving micro-arcsec structure in the lensed cosmologically-distant source if it crosses a caustic created by the stellar-mass compact objects in the lens mass distribution (e.g. Chang \\& Refsdal 1984; Grieger, Kayser \\& Refsdal 1988; Wambsganss, Paczy\\'nski \\& Schneider 1990; Wo\\'zniak et al. 2000). Moreover, because of the presence of multiple images, one can in principle separate intrinsic source fluctuations from microlensing variability.  For example, ongoing microlensing of the lensed (optical) images in the system Q2237+0305 (Huchra et al. 1985) has been observed ever since its discovery (e.g. Irwin et al. 1989; Corrigan et al. 1991; {\\O}stensen et al. 1996; Lewis et al. 1998; Wo{\\'z}niak et al. 2000). To a smaller degree and on longer time scales, microlensing in Q0957+561 has been detected as well (Pelt et al. 1998; see also Refsdal et al. 2000). The time scale of microlensing in both cases is defined by the relative transverse velocity between the source, lens and observer, which is given by the bulk velocity of the lensing galaxy plus the random motions of compact objects in the line-of-sight to the stationary quasar images, and is typically of order several hundred km\\,s$^{-1}$. This results in microlensing time scales of the order of months to years for solar-mass objects (e.g. Wambsganss 2000).  \\begin{figure*} \\begin{center} \\leavevmode \\hbox{% \\epsfxsize=\\hsize \\epsffile{fig1r.eps}} \\end{center} \\caption{A cartoon of how we determine the microlensing magnification as a function of time for the expanding shell source superimposed on the magnification patterns caused by massive compact objects. The five panels indicate five different epochs (in practise, we consider 100 epochs). The superimposed curved line indicates the microlensing magnification as a function of time, corresponding to this particular example. The arrow points to a particular high magnification structure that is passed by the ring source at epoch $t_2$.} \\end{figure*}   Besides these optical sources, recently the first case of radio-microlensing was reported in the lens system B1600+434 (Koopmans et al. 2000a, 2000b; Koopmans \\& de Bruyn 2000), suggesting extremely compact relativistic substructure in the lensed radio source. If this microarcsec-scale substructure is part of a relativistic jet, the time-scale of microlensing variability by stellar-mass objects reduces to several weeks (Koopmans \\& de Bruyn 2000), allowing one to probe compact objects up to $\\sim$$10^{5}$\\,M$_\\odot$ on time scales of several years, as well as the substructure of the relativistic jet. Indications of optical microlensing in B1600+434 with similar time scales have also been found (Burud et al. 2000). Clearly, microlensing is a promising field of future research regarding the study of high-$z$ sources at micro-arcsecond scales.  In addition to testing the structure of AGNs, Loeb \\& Perna (1998) more recently proposed the use of microlensing to probe into the internal structure of GRB afterglows on micro-arcsec scales. Garnavich, Loeb \\& Stanek (2000) indeed interpreted an anomalous event in the lightcurves of the optical afterglow of GRB000301C (Masetti et al. 2000; Sagar et al. 2000; Berger et al. 2000; Jensen et al. 2000; Smette et al. 2000) as being caused by microlensing of the GRB afterglow. They subsequently derive constraints on its structure, which appear in good agreement with theoretical blastwave models. The inferred mass of the lensing object is $\\sim$0.5\\,M$_\\odot$, if its redshift is optimal for microlensing (i.e. about half way; Garnavich et al. 2000).  However, one needs to be cautious here, because no multiple images are present -- as in the case of strong gravitational lenses -- to confirm that this is indeed a non-intrinsic event. The fact that the event occurs within only a few days after the burst and has an amplitude of $\\approx$1 mag, suggests that the burst must have occurred close to the Einstein radius of an intervening massive compact object (Garnavich et al. 2000). To have a significant probability of observing such a GRB microlensing event, the universe requires a surface density in compact objects close to the critical surface density (e.g. Press \\& Gunn 1973; Blaes \\& Webster 1992).  The situation is partly similar to the case of variability in single quasars (i.e. not multiply-imaged), which also cannot easily be proven to be due to microlensing (e.g. Hawkins \\& Taylor 1997; Hawkins 1998), although the freedom to model the GRB afterglow lightcurve (mostly dominated by self-similar expansion resulting in a power-law behavior) is significantly less than that for quasars.  In this paper, we investigate the probability of microlensing in GRB afterglows for a flat $\\Lambda$--dominated cosmological model and with the distribution of massive compact objects connected to visible galaxies, particularly focusing on the strong event seen in the GRB000301C afterglow. Our conclusions, however, are independent of this particular GRB. In Section 2, we describe numerical microlensing simulations of GRB afterglow light curves, extended to high microlensing optical depths. In Section 3, the {\\sl a posteriori} probability of the observed event in GRB000301C is calculated. Section 4 summarizes our result and states our conclusions. ",
        "conclusions": "If the event seen in the optical lightcurve of the GRB000301C afterglow is caused by microlensing, the opportunities to study the structure and evolution of GRB afterglows on micro-arcsec scale are potentially very exciting (Loeb \\& Perna 1998; Garnavich et al. 2000; Mao \\& Loeb 2000).  However, in this paper we have shown that the {\\sl a posteriori} probability of this particular event is actually small, i.e. 0.7--2.7\\% for a fraction $f_*$=0.2--1.0, respectively, of dark-matter in the form of compact objects. The main assumptions in this calculation are: (i) a constant comoving density population of galaxies which follow a Schechter luminosity (i.e. mass) function, (ii) a flat $\\Lambda$--dominated FRW universe with $\\Omega_{\\rm m}$=0.3, (iii) all matter in the universe traces galaxies, which can be described as singular isothermal spheres, (iv) the mass spectrum of compact objects is `narrow', in which case they can be parameterised by a single value for their mass, (v) a typical GRB optical afterglow has similar properties as GRB000301C and has a redshift of $z$$\\approx$2, and (vi) there is no significant magnification bias. The magnification bias, however, is unlikely to be a problem. Even in case this bias increases the number of observed GRBs by a factor of 10 for high microlensing optical depths ($\\kappa$$\\ga$0.25), Table~1 and eqn.(7) show that the overall probabilities are increased by only 10$\\times P(\\Delta m \\ge 0.95^{\\rm m}|\\kappa\\ga0.25) \\times\\tau_{\\rm GL/GRB}$$\\sim$0.1\\%. A detailed calculation should consistently take into account both the magnification distribution of galaxies {\\sl and} that of compact objects (e.g. Pei 1993a, 1993b). In light of the very uncertain redshift distribution and luminosity function of GRB afterglows this is, however, not yet warranted.  Moreover, even in the case that the {\\sl whole} sky has a surface mass density of $\\kappa_*$$\\la$0.25 (normalised by the ``critical'' surface mass density for lensing, e.g. Schneider et al. 1992) our simulations (see Table~1) indicate that the probability of such an event does not significantly exceed $\\approx$5\\%.  Clearly, the latter situation is unrealistic, but it does place a very robust upper limit on the probability of this event (see also Press \\& Gunn 1973 and Blaes \\& Webster 1992).  The plausibility to actually observe such a fraction of GRB afterglow lightcurves with a microlensing event of about one magnitude also depends on the typical mass of compact objects.  However, even if the mass goes up or down by a factor of ten from the assumed 0.5\\,M$_\\odot$, to first approximation that would just mean a ``broadening'' in time of the microlensing event by about a factor of three (since the length/time scale is proportional to the square root of the mass). However, if most compact objects have masses {\\it much} smaller than that inferred for GRB000301C, the fraction of GRB afterglows with microlensing signatures will obviously decrease. In that case the effect of the ``shrunk'' caustics will average out over the much larger angular size of the emission region of the GRB afterglow. This case, however, seems unlikely in light of the stringent lower limit ($\\sim$10$^{-2}$\\,M$_\\odot$) placed on low-mass compact objects in 0957+561 (Schmidt \\& Wambsganss 1998; Refsdal et al. 2000; Wambsganss et al. 2000). The fraction of microlensed GRB afterglows would also decrease if the typical redshift of GRBs is significantly lower than~$z=2$. A somewhat lower typical redshift is supported by the average redshift $<z>_{\\rm GRB}\\approx1.3$, with an rms spread of $\\approx$1, that we find from Bloom, Kulkarni, \\& Djorgovski (2000). In case of a broad mass spectrum of the compact objects or a large spread in the physical properties of the GRB optical afterglows, one has to convolve the microlensing probabilities with these properties. This, however, is not expected to change the results significantly.  We emphasise that our results {\\sl do not exclude} that the event seen in the lightcurve of the optical afterglow of GRB000301C is in fact due to gravitational microlensing, especially in light of the fact that a galaxy was found at 2~arcsec from the GRB (Garnavich et al. 2000). Each event should be treated on its own merit (and it is a posteriori anyway). In any case, in the sample of $\\sim$20 known GRBs with observed optical afterglows (e.g. Bloom et al. 2000), the probability of one such event would be around 14\\%.  But our results do predict that it is unlikely to find many strongly microlensed GRB afterglow soon, and it emphasises the difficulties one will encounter in creating a large sample of these strong (i.e. $\\Delta m \\ga1$\\,mag) events, with the aim to statistically study the physical properties of GRB afterglows on micro-arcsec scales.  For example, to get 10 additional microlensing events of comparable strength, one would require $\\approx$1500 GRB afterglow optical lightcurves. Even for the {\\sl Swift} satellite (Parsons et al. 1999; see also Mao \\& Loeb 2000) this provides a challenging task.  On the bright sight, if the event in GRB000301C was due to microlensing, we can expect a significant fraction of high-$z$ GRB afterglow light curves to show microlensing events stronger than 0.1\\,mag (see also Mao \\& Loeb 2000), although these will be hard to distinguish from intrinsic variations and will also not really be able probe the GRB afterglow structure on micro-arcsec scales. We conclude, that the best cases to unambiguously study micro-arcsec structure in GRB afterglows will probably be those GRBs that are multiply imaged (about 1--2\\% expected, see, e.g.  Holz et al. 1999; Komberg et al. 1999; Mao 1992, 1993; Marani et al. 1999; McBreen et al. 1993; Narayan \\& Wallington 1992; Nowak \\& Grossman 1993; Paczy\\'nski 1987; Wambsganss 1993; Williams \\& Wijers 1997).  Those (soon to be discovered!) will almost certainly (see Table~1) show microlensing events $\\ga$0.1\\,mag, which can easily be separated from intrinsic variations by comparison with the other lensed GRB images, after correcting for the respective time delays. A smaller fraction ($\\approx$5\\%) will also show events $\\ga$1\\,mag, although this might increase due to the magnification bias. So the prospects of doing science with (micro-)lensed gamma ray bursts and afterglows remain promising."
    },
    "0011/astro-ph0011503_arXiv.txt": {
        "abstract": "Hydrodynamical winds from a spherical two-temperature plasma surrounding a compact object are constructed. The mass-loss rate is computed as a function of electron temperature, optical depth and luminosity of the sphere, the values of which can be constrained by the fitting of the spectral energy distributions for known X-ray binary systems. The sensitive dependence of the mass loss rate with these parameters leads to the identification of two distinct regions in the parameter space separating wind-dominated from non wind dominated systems. A critical optical depth ($\\tau_c$), as a function of luminosity and electron temperature, is defined which differentiates these two regions. Systems with optical depths significantly smaller than $\\tau_c$ are wind-dominated.  The results are applied to black hole candidate X-ray binary systems in the hard spectral state (Cyg X-1, GX 339-4 and Nova Muscae), and it is found that the inferred optical depth ($\\tau$) is similar to $\\tau_c$ suggesting that they are wind regulated systems. On the other hand, for X-ray binary systems containing a neutron star (e.g., Cyg X-2) $\\tau$ is much larger than $\\tau_c$ indicating the absence of significant hydrodynamical winds. ",
        "introduction": "\\label{sect1}  X-ray binary systems containing a black hole candidate are typically found to be in two different spectral states. In the hard state, the broad band X-ray spectrum can be described as a power-law (photon spectra index $\\approx 1.5$) with a high energy cutoff around 100 keV. In the soft state, the spectrum consists of two components. There is usually an extended power-law (with spectral index $\\approx 2.5$) and a soft X-ray emission which has a spectral shape similar to that of a black body. For a recent review of the observations and phenomenological description of these sources see Tanaka \\& Shibazaki (1996).  The modeling of the hard state spectra can be described in terms of an unsaturated Comptonization process of soft photons in a region with hot electrons ($kT \\approx 50$ keV) and electron scattering optical depths of order unity. In a pioneering study, Shapiro, Lightman \\& Eardley (1976) identified this hot region with a geometrically thick, optically thin, hot accretion disc. In this model, the gravitational energy dissipated in the disc, heats the ions which in turn transfer the energy to electrons by Coulomb interactions. However, the electron-ion Coulomb interaction rate is inefficient in such an environment of low density and high electron temperature, and this leads to a large difference in temperature between the electrons and ions, with the ion temperature reaching nearly virial values ($10^{11} K$). The importance of radial advection of energy in such a disc was noted by Ichimaru (1977). Taking advection into account Narayan \\& Yi (1994) constructed self-similar solutions for the disc equations called Advection Dominated Accreting Flows (ADAF) showing not only that the proton temperatures approach their virial values, but also that the radiative efficiency of accretion can be significantly reduced as a result of the advection of energy into the black hole. In an alternative description Chakrabarti \\& Titarchuk (1995) argued that under certain conditions a shock may arise in such accretion discs and identified the hot Comptonizing region with the post-shock flow. Despite the differences in the geometry, radiative processes and the detailed disc structure, both these models have in common the presence of a two-temperature plasma.  Such a plasma is a natural outcome of any accretion disc model which (a) identifies the hard state X-ray spectrum as a result of the unsaturated Comptonization process of soft photons, (b) assumes that the viscous energy dissipated heats the ions and (c) that the only mechanism for energy transfer between the ions and electrons is Coulomb interaction. We note here that these assumptions may not be valid since the viscous energy dissipated may heat the electrons preferentially if a strong equipartition magnetic field is present in the disc (Bisnovatyi-Kogan \\& Lovelace 2000).  In addition, there could be unknown mechanisms which transfer energy between ions and electrons more efficiently than the Coulomb interaction. Thus, it will be useful to have an independent observational signature which could confirm the existence of two-temperature plasmas in black hole candidate systems.  The nearly virial  proton temperature of this plasma suggests the possibility of a strong hydrodynamical wind arising from these systems (e.g., Piran 1977; Takahara, Rosner \\& Kusunose 1989; Kusonose 1991). Such an outflow could transport away a significant fraction of mass, energy and/or angular momentum, thereby affecting the structure and stability of the disc and affecting the radiative efficiency of accretion for a given mass transfer rate. Chakrabarti (1999) and Das (1999) have studied the possibility of outflows in the context of the shock/centrifugal barrier models. They find that for certain values of the disc parameters (e.g., accretion rate and specific entropy) a hydrodynamical wind occurs. For the ADAF disc solutions, Blandford \\& Begelman (1999) argued that only a small fraction ($<$ 1\\%) of the gas actually falls into the black hole and the rest is driven away as a wind (in an advection dominated inflow-outflow solution - ADIOS), thereby effectively reducing the radiative efficiency. Beckert (2000) confirmed this result for different viscosity laws while Quataert \\& Narayan (1999) showed that the X-ray spectra from such a wind driven accretion process can explain the observed spectra of some black hole candidate systems in quiescence. These calculations were undertaken for low mass accretion rates, and it is not clear how the system will behave in the high accretion rate regime inferred for black hole candidate systems in the hard state.  The formation of hydrodynamical winds from accretion discs depends on the structure and geometry of the discs. Thus, detailed calculations of the outflow are intrinsically model dependent. Further, reliable calculations of the structure of such discs are difficult because uncertainties exist in the vertical distribution of energy dissipation in the disc associated with our lack of detailed understanding of the viscosity.  We note that internal magnetic fields in the disc could also facilitate (or inhibit) the formation of winds. Here, electromagnetic forces may accelerate and collimate the wind to form high velocity jets as observed in microquasar systems (Mirabel \\& Rodriguez 1999).  Considering these uncertainties a prudent approach would be to estimate the wind characteristics using only those parameters which can be directly constrained by the fitting of spectral energy distributions. Such an analysis would allow a rough estimation of the magnitude of the mass and energy lost in the form of a wind  for a system in question. With such an objective in mind, we report in this paper on calculations of the structure of an hydrodynamical wind similar to those found by Takahara et al (1989), but in the context of a uniform spherical two-temperature plasma around an accreting black hole. Rather than treating the entire disc/wind configuration with detailed heating and cooling processes as in Kusunose (1991) the calculations are parameterized in terms of the electron temperature of the cloud, $T_e$, the optical depth, $\\tau$, and the total luminosity of the source.  We adopt this approach as these parameters are constrained by spectral fitting analyses. In the next section, the formulation of the problem and the numerical results are presented.  The application of these results to observed systems is given in \\S 3 and discussed in the last section. ",
        "conclusions": "\\label{sect4}  The possible occurrence of hydrodynamical winds in X-ray binary systems has been investigated.  From simple considerations, the conditions under which these winds can be important have been identified.  It is found that the mass loss rate in such winds depends sensitively on the luminosity of the source as well as the electron temperature and optical depth of the coronal region.  The steep dependence of the mass loss rate on these parameters facilitates the use of a critical optical depth to indicate whether a given system can support such a wind. Application of the theory indicates that winds can exist in the hard state of the black hole candidates Cyg X-1, GX 339-4 and GS 1124-68.  Strong winds in these systems may decrease the radiative efficiency ($\\eta = L/\\dot M_i c^2$, where $\\dot M_i$ is the mass inflow rate) by carrying away a substantial amount of matter and energy. For Nova Muscae, the radiative efficiency has been estimated to be $\\eta = 0.01$ for the soft state and $\\eta = 0.05$ for the hard state (Misra 1999). These values are lower than that expected from an Keplerian disc ( $\\eta \\approx 0.1$) which could be due to the presence of a strong wind rather than energy advection into the black hole. Winds may also effect the thermal stability of accretion discs by introducing an additional channel for energy loss (e.g., Piran 1977).  We reiterate that the analysis undertaken in this paper is based on the assumption that the gravitational energy dissipated in the system is transferred to the electrons by the ions via Coulomb interactions. This naturally restricts the analysis to only certain systems where this is valid. As mentioned earlier, the soft state spectrum of black hole candidate systems is probably of a non-thermal origin and, hence, the analysis performed in this study cannot be applied. Even for the hard spectral states of compact X-ray binaries considered here, alternate models to a two -temperature plasma description have been proposed. For example, the X-ray spectra could be due to magnetic flare activity above a cold disc (Poutanen \\& Fabian 1999) or produced by a disc with a rapidly varying radial temperature profile (Misra, Chitnis  \\& Melia 1998). Thus, it should be emphasized that the results presented here are specifically discussed within the framework of a two-temperature plasma model.  The spherical geometry assumed here is simplistic even though two-temperature discs are geometrically thick. Further the effect of angular momentum or of convection on the dynamics of accretion and the wind outflow have not been taken into account. The former effect can increase the mass outflow rate as a result of centrifugal support, whereas the latter effect can affect the thermal structure of the underlying disk and, hence, the existence of a wind.  In this context, the thermal structure is dependent on the direction in which angular momentum is radially transported by convection and the magnitude of the viscosity parameter (see Narayan, Igumenshchev, \\& Abramowicz 2000). A more detailed analysis should take these effects and radiative heating/cooling of the wind into account in at least two spatial dimensions (e.g., Chakrabarti \\& Molteni 1993). However, such analyses will, by necessity, be model dependent and limited by the uncertainties in the vertical structure and geometry of the hot disc.  The hydrodynamical winds described in this paper may be confined to form a jet-like structure by the geometry of the disc and magnetic fields. Further if they are accelerated to relativistic speeds by electromagnetic forces, they may provide the origin of the radio jets observed in the black hole candidate systems known as microquasars (Mirabel \\& Rodriguez 1998).  On the other hand, the radio jets may be a different phenomenon unrelated to hydrodynamical winds. In that case it is desirable to have direct observational signatures of these outflows. Since the Thomson optical depth of the winds calculated here (i.e. the optical depth from the surface of the sphere $R$ to infinity) is typically $< 0.1$, column densities of the order of $10^{23}$ cm$^{-2}$ are indicated. Although the X-ray continuum spectra are not expected to be altered by the outflow, a significant fraction of the column density may not be highly ionised giving rise to observable absorption and/or emission lines. The detection of P-Cygni type profiles by high resolution X-ray satellites Chandra and Newton-XMM may provide evidence for the existence of such winds and serve as a useful diagnostic not only of the wind structure, but also the physics underlying the disc as well.  \\bigskip RM acknowledges support from the Lindheimer Fellowship at Northwestern University."
    },
    "0011/astro-ph0011359_arXiv.txt": {
        "abstract": "We discuss the ultraviolet to near-IR galaxy counts from the deepest imaging surveys, including the northern and southern {\\it Hubble Deep Fields}. The logarithmic slope of the galaxy number-magnitude relation is flatter than $0.4$ in all seven $UBVIJHK$ optical bandpasses at faint magnitudes, i.e. the light from resolved galaxies has converged from the UV to the near-IR. Most of the galaxy contribution to the extragalactic background light (EBL) comes from relatively bright, low-redshift objects (50\\% at $V_{\\rm AB}\\lta 21$ and 90\\% at $V_{\\rm AB}\\lta 25.5$). We find a lower limit to the surface brightness of the optical EBL of about $15\\,\\eblunits$, comparable to the intensity of the far-IR background from {\\it COBE} data. Diffuse light, lost because of surface brightness selection effects, may add substantially to the EBL. ",
        "introduction": "The extragalactic background light (EBL) is an indicator of the total luminosity of cosmic structures, as the cumulative emission from pregalactic, protogalactic, and evolving galactic systems, together with active galactic nuclei (AGNs), is expected to be recorded in this background. The recent progress in our understanding of faint galaxy data, made possible by the combination of {\\it Hubble Space Telescope} {(\\it HST}) deep imaging and ground-based spectroscopy, and of the evolution of the stellar birthrate in optically-selected galaxies from the present-epoch up to $z\\approx 4$ (Steidel \\etal 1999; Madau, Pozzetti, \\& Dickinson 1998), has been complemented by measurements of the far-IR/sub-mm background by the {\\it COBE} satellite (Hauser \\etal 1998; Fixsen \\etal 1998; Puget \\etal 1996), showing that a significant fraction of the energy released by stellar nucleosynthesis is re-emitted as thermal radiation by dust (Dwek \\etal 1998).  In this talk I will focus on the galaxy number-apparent magnitude relation and its first moment, the galaxy contribution to the EBL. The logarithmic slope of the differential galaxy counts ($d\\log N/dm_{\\rm AB}\\equiv \\gamma$) is a remarkably simple cosmological probe of the history of the stellar birthrate, as it must drop below 0.4 to yield a finite value for the EBL. The radiation emitted from unresolved sources that could be lost due to uncertainties in the faintest galaxy data and surface brightness selection effects will be discussed, together with the contribution to the EBL from high-$z$ populations such as the Lyman-break galaxies and extremely red objects.  \\begin{figure} \\centerline{\\psfig{figure=EBL_diff.ps,width=8cm}} \\caption{\\small Extragalactic background light per magnitude bin, $i_\\nu=10^{-0.4(m_{\\rm AB}+48.6)}N(m)$, as a function of $U$ ({\\it filled circles}), $B$ ({\\it open circles}), $V$ ({\\it filled pentagons}), $I$ ({\\it open squares}), $J$ ({\\it filled triangles}), $H$ ({\\it open triangles}), and $K$ ({\\it filled squares}) magnitudes. For clarity, the $BVIJHK$ measurements have been multiplied by a factor of 2, 6, 15, 50, 150, and 600, respectively.} \\end{figure}  \\begin{figure} \\centerline{\\psfig{figure=EBL_cum.ps,width=7cm}} \\caption{\\small Cumulate contribution to the EBL per magnitude bin as a function of $U,B,V,I,J,H,K$ AB magnitudes, derived from a fit to the observed counts.} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011390_arXiv.txt": {
        "abstract": "s{ We present the results of a redshift survey of both cluster and intercluster galaxies in the central part of the Shapley Concentration, the richest supercluster of clusters in the nearby Universe, consisting of $\\sim 2000$ radial velocities. We estimate the total overdensity in galaxies of the supercluster and its mass and we discuss the cosmological implications of these results. Moreover, using a Principal Components Analysis technique, we study the influence of the cluster and supercluster dynamics on the galaxy spectral morphology. }  \\noindent Clusters of galaxies are thought to form by accretion of subunits in a hierarchical bottom up scenario: numerical simulations revealed that mergings happen along preferential directions, called density caustics, which define matter flows, at whose intersection rich clusters are formed. \\\\ The central regions of rich superclusters, where the overdensity is of the order of $\\sim 10$, are thought to be in the early collapse phase and therefore give the possibility to study the formation of clusters and groups of galaxies. These regions are the ideal environment for the detection of cluster mergings, because the peculiar velocities induced by the enhanced local density of the large scale structure favour the cluster-cluster and cluster-group collisions, in the same way as the caustics seen in the simulations. \\\\ Spectacular major merging events are visible in the central region of the Shapley Concentration, the richest supercluster of galaxy clusters in the nearby Universe (Zucca et al. 1993). Particularly interesting are two chains of interacting clusters: the A3558 and the A3528 complexes (Fig. 1). \\\\ We are carrying on a long term multiwavelength study on the Shapley Concentration (see also the poster by Bardelli et al., these proceedings), in order to determine its mass, dynamical state and cosmological relevance, as well as to analyse in detail the physics of the merging phenomena. \\\\ In this poster we focus on the large-scale distribution and spectral properties of galaxies both in clusters and in the ``field\" in between clusters.  \\begin{figure} \\begin{center} \\epsfxsize=0.7\\hsize \\rotate[r] { \\epsfbox {fig1a_iap.ps} } \\epsfxsize=0.7\\hsize \\rotate[r] { \\epsfbox {fig1b_iap.ps} } \\end{center} Figure 1: Optical distribution of galaxies in the central part of the Shapley Concentration. \\end{figure} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011445_arXiv.txt": {
        "abstract": "The time evolution of initially balanced, rapidly rotating models for an isolated disk of highly flattened galaxies of stars is calculated. The method of direct integration of the Newtonian equations of motion of stars over a time span of many galactic rotations is applied.  Use of modern concurrent supercomputers has enabled us to make long simulation runs using a sufficiently large number of particles $N=30,000$.  One of the goals of the present simulation is to test the validities of a modified local criterion for stability of Jeans-type gravity perturbations (e.g. those produced by a barlike structure, a spontaneous perturbation and/or a companion galaxy) in a self-gravitating, infinitesimally thin and collisionless disk. In addition to the local criterion we are interested in how model particles diffuse in velocity.  This is of considerable interest in the kinetic theory of stellar disks.  Certain astronomical implications of the simulations to actual disk-shaped (i.e. rapidly rotating) galaxies are explored.  The weakly nonlinear, or quasi-linear kinetic theory of the Jeans instability in disk galaxies of stars is described as well. ",
        "introduction": "In galactic astronomy, the most intensely studied of the problems of pattern formation is the problem of the formation of spiral structure, which has brought forth numerous theoretical and numerical approaches. Even though at the present time there exists, as yet, no satisfactory theory of the origin and conservation of the spiral structure so prominent in the Milky Way and many other giant flattened galaxies, the majority of experts in the field has yielded to the opinion that the study of the stability of the small-amplitude cooperative oscillations in disk-shaped galaxies of stars is the first step towards an understanding of the phenomena.  This is because in the Milky Way and many other disk galaxies, the bulk of the luminous mass, probably $\\gtrsim 90\\%$, is in the stars.  Therefore, it seems likely that the spiral structures are intimately connected with the stellar constituent of a galaxy, and that stellar dynamical phenomena play a basic role.  Because of the long-range nature of the gravitational force between stars, a galaxy exhibits collective modes of motions --- modes in which the stars in large regions move coherently or in unison.  A galaxy is characterized also by a high ``temperature\" (a dispersion of random velocities of stars) and hence a high ``thermal\" energy; this thermal energy is much larger than the interaction potential between pairs of stars. Because of this, binary encounters produce only small-scale scattering of the motions of the stars.  Over long enough times, these gradually build up to produce large deflections that constitute collisions.  In many cases these collisional effects are so weak that we consider the galaxy as collisionless on the Hubble time $T \\sim 10^{10}$ yr (Chandrasekhar, 1960; Binney and Tremaine, 1987).  Thus, the galaxy is dominated by the collective motions and the free streaming of the stars (kinetic effects).  As usual it is very difficult determine a priori whether a particular linearization of nonlinear equations made in analytical studies constitutes a valid approximation.  This can be determined, however, by direct computer simulations of the nonlinear theory that mimic the behaviour of stars in galaxies. Such modeling gives more detailed information than can be obtained analytically, so that the important physical phenomena can be determined.  In this paper, the results of the approximate theoretical analysis described in the Appendix are compared with our own many-body ($N$-body) simulations.   One of the goals of the simulation is to test the validities of a modified local criterion for stability of Jeans-type gravity perturbations in a self-gravitating, infinitesimally thin and collisionless disk.  The fact that the nonaxisymmetric perturbations in the differentially rotating system are more unstable than the axisymmetric ones is taken into account in this modified Toomre-like (Toomre, 1964, 1977) criterion.  In addition to the stability criterion we are interested in how model particles diffuse in velocity.  As for the present study, such random velocity diffusion is caused by stellar disk turbulence. In particular, we verify the analytical prediction that as the Jeans-unstable perturbations grow the mean-square velocity increases linearly with time.  The observation of this behavior is a sensitive test of the model as well as the kinetic theory of stellar disks.  In fact, Morozov (1981) and Griv {\\it et al.} (1999a) already attempted to confirm the modified criterion numerically.  However, because of the very small number of model stars, $N=200$, Morozov's results are subject to considerable uncertainties, and additional simulations are clearly required to settle the issue.   Furthermore, for that number of model particles, the two-body relaxation time scale is comparable to the crossing time, even with Morozov's modest softening parameter, raising some question about the applicability of his simulations to actual almost collisionless galaxies.  Increasing the number of model stars is definitely a more reliable procedure.  In turn, Griv {\\it et al.} used a different numerical approach of the so-called local $N$-body simulations.  The $N$-body experiments in a local or Hill's approximation has been pioneered by Toomre (1990), Toomre and Kalnajs (1991) and Salo (1995). In these simulations dynamics of particles in small regions of the disk are assumed to be statistically independent of dynamics of particles in other regions. The local numerical model thus simulates only a small part of the system and more distant parts are represented as copies of the simulated region.  Unlike Morozov (1981) and Griv {\\it et al.} (1999a), we study some aspects of dynamical behavior of stellar systems by global simulations using a sufficiently large number of particles. ",
        "conclusions": "The results show that a velocity dispersion given by Toomre's (1964) criterion will stabilize a disk only against axisymmetric ring-like gravity perturbations.  However, such disks are unstable against nonaxisymmetric spiral-like perturbations.  Jeans-unstable spiral perturbations ``heat\" the system; the mean-square random velocity increases linearly with time.  The results are in agreement with analytical predictions.  We were able to generate an axisymmetric, gravitationally stable disk. The initial condition for the axisymmetric stable disk was obtained analytically and confirmed experimentally: the radial dispersion of random velocities of stars should be equal (or greater than) to the modified dispersion at each radii."
    },
    "0011/hep-ph0011310_arXiv.txt": {
        "abstract": " ",
        "introduction": "Many ultra high energy (UHE) cosmic ray air showers with energies in excess of $5\\times 10^{19}$ eV have been observed in the past few decades \\cite{nw}. The nature and origin of the primary particles is not understood \\cite{nw,es,sigl}.  The puzzle is that the sources have to be within the GZK limit of approximately 50 Mpc if these particles are protons or nuclei \\cite{GZK,puget}. However there are not enough powerful astrophysical sources within this distance to explain the events.  Among known particles, only neutrinos travel larger distances than protons in intergalactic space. This leaves neutrinos as the only established candidates that can travel the  distances greater than 100 Mpc from known UHE sources.  The GZK bound of 50 Mpc is not applicable to them. Yet neutrino interactions with matter are too weak in the Standard Model of particle physics to generate the observed air showers. Hence these events seem to demand a revision of our current understanding of nature. Either the determination of the number of sources of such ultrahigh energy particles in our astrophysical neighborhood is grossly low\\footnote{For example if magnetic fields outside the galaxy have been underestimated, ``line of sight'' and ``photon travel time'' requirements on protons and nuclei can be relaxed and new source possibilities considered \\cite{es}}, or these observations are a signal of new physics.  Many speculative ideas have been proposed to explain the events above $10^{19} - 10^{20}eV$, including topological defects such as cosmic strings \\cite{abs,mb} and associated decays of heavy, relic particles \\cite{berez,kt}, existence of neutral, stable, strongly interacting particles, such as a light gluino \\cite{ch,raby,lou} or a monopole bound state \\cite{kw,wick}, and violation of Lorentz invariance \\cite{g-m,cg}.  Much of this work requires that the primary particle responsible for generating these air showers is an exotic new particle which does not exist within the Standard Model.  In a recent paper \\cite{us1} we argued that the data is consistent with the general features of massive spin-2 exchange.  Models where effects of gravity can be strong just above the weak scale \\cite{add} supply a natural and attractive framework.  The interaction cross section of neutrinos with matter is greatly enhanced with massive spin-2 exchange at UHE and may reach values close to the hadronic cross sections. In the low scale gravity models, the cross section enhancement arises from t-channel exchange of the tower of gravitons.  Our estimates of the neutrino-proton cross section at the highest energies relevant for these events are of the order of one to a few hundred millibarns. The highest energy cosmic ray events may therefore be initiated by neutrinos\\footnote{Correlation between the positions of compact radio quasars and the track directions of UHE $>100$ EeV cosmic rays has been studied by several groups \\cite{fb,sigl2,us2}.}\\cite{us1,early}.  A generic, robust prediction of massive spin-2 exchange, known for many decades, is that the total cross section should grow with a power of energy, typically $\\sigma_{tot} \\sim s^2$. The property of power law growth with a power exceeding 1 (the result of 4-Fermi {\\it spin-1} exchange) is quite hard to evade and can be traced to dimensional analysis.  We consider large cross sections at $UHE$ to be characteristic of extra-dimension, low scale gravity models. Interaction of $UHE$ neutrinos is a quite natural domain to seek the new effects of low scale gravity models:  the very weakness of the Standard Model neutrino coupling minimizes this background, while the regime of of highest possible energy maximizes the effects of graviton-KK mode exchange.  The theory of low-scale gravity models is only partly developed, and questions of unitarity complicate the interpretation of perturbation theory \\cite{gw}.  One can choose models of the cross section which are further from the calculations of perturbation theory in the sense that they grow at a slower rate with energy than $s^2$ (The perturbative, parton level cross section rises as $\\hat s^3$) \\cite{N+S,kp,usnote}, or which operate by a separate (s-channel) mechanism \\cite{d+d}.  Indeed it is possible to restrict models of low-scale gravity to the extent that nothing observable is predicted at the energies in question. For example, the astrophysical bound on the scale parameter $M$ for the n=2 case guarantees that the consequences of this model are unobservable \\cite{N+S},\\cite{kp}.  There has been some confusion on this point.  Let us compare the model we use here, taken from our previous cross section calculations \\cite{us1}, with subsequent work \\cite{kp}. The latter reports the result of assuming that a finite brane tension introduces an exponential damping of higher KK modes, providing an alternative cutoff mechanism \\cite{bando}.  Like our calculation, when $\\surd s \\geq M$, the cross section in \\cite{kp} rises approximately quadratically with neutrino energy (See Fig. 1 in \\cite{kp}, where $\\sigma_{\\nu N}$ rises by two orders of magnitude for every order of magnitude rise in $E_{\\nu}$).  Unlike ours, the calculation there assumes n=2 only, for which SN1987a analysis makes the restriction $M$ $\\geq$ 30-70 TeV \\cite{bounds}.\\footnote {Citing uncertainties in astrophysical parameters, \\cite{kp} considers $M$ values as low as 6 TeV.}  If $n \\geq 3$ were considered, the scale could be lowered to the 2-3 TeV range and cross section values in agreement with ours would result.  This is clear from the trend with mass scale in Fig. 1 in \\cite{kp}.  Conversely, we could suppress our cross section to their values by raising $M$ to values of 6 TeV and above.  Specifically, we find that $\\beta = 1$ and $M$ = 6 TeV yields $\\sigma_{\\nu N}$ = 0.3 mb, compared to 0.1 mb at $E_{\\nu}$ = $10^{20}$ TeV for $M$ = 6 TeV in \\cite{kp}, while the choice $\\beta$ = 2 and $M$ = 6.6 TeV or $\\beta$ = 1 and $M$ = 7.3 TeV reproduces their 0.1 mb value.  Within modest parameter variations, the results clearly agree.  This is not a surprise, since the parton level amplitudes and cross sections for small t are essentially identical, and are insensitive to the value of n in the two cutoff methods \\cite{bando}. The two calculations differ only in the details of the large t cutoffs, both of which produce $s^{2}$ behavior of the cross sections.  The cutoff used in in \\cite{kp} gives a cross section result in essential agreement with ours at a given set of $E_{\\nu}$ and $M$ values. Their assertion to the contrary is an unfortunate consequence of presenting the results for a lower bound $M \\geq 6$, justified only for n=2, and drawing sweeping, unjustified conclusions about the general situation. With this technical issue clarified about the particular model we employ, we reiterate that our goal is to explore the broad consequences of strongly interacting $UHE$ neutrinos.  The primary, general question that arises, defining our focus here, is the nature of {\\it  air shower development}. High-energy leptonic interactions do not have the same multiplicity or inelasticity as high energy hadronic interactions.  To understand the potential relevance of neutrino interactions, one must address not only the total cross section, but also the way the interaction delivers energy into the air showers that are actually observed.  In the present paper we compare simulations of air showers generated by neutrino primaries with large cross sections to those generated by protons in the Standard Model.  We ask whether there is anything about existing showers which might {\\it rule out} neutrinos as primaries.\\footnote{Alternatively one might ask whether one can ``find evidence'' for neutrinos as the primaries in some features of the showers.  We do not pursue this, because the fluctuations of air showers and flexibility in simulation codes make it a very hard and ambiguous way to proceed.} If so then the case is made that large cross sections alone are not enough to support the case for neutrinos, and speculative models of \\emph{new} particles might be indicated.  Contrary to some expectations \\cite{kp}, we find that neutrinos with large cross sections {\\it can create showers that are much like proton-initiated showers} and that in some cases are indistinguishable from them \\cite{sm} \\cite{ddm}. Two features of the low-scale gravity contribution to the neutrino-nucleon cross section come into play: first, the cross section is large enough to initiate air showers at high enough altitudes; second, its rapid $s^2$ dependence suppresses new effects among secondary products, which carry at most a few percent of the primary energy.  Our  methodology can evidently be extended to other models for hadronic size UHE neutrino cross sections provided the cross sections grow rapidly (as in spin-2 exchange).  Other speculative primaries should be considered on a case-by- case basis.  One cannot take the existence of one model that produces well simulated showers above $10^{20}$ eV to be conclusive evidence for a given hypothesis for the identity of the primary, neutrino or otherwise.  The question of the mysterious primary, then, needs to be framed in view of everything that can be observed:  cross sections, shower characteristics, and angular distributions and correlations, which may be informative about the charge of the primary. ",
        "conclusions": "The conjecture that UHE, super-GZK cosmic ray showers are caused by neutrinos with large UHE cross sections is almost as old as the field itself \\cite{early}.  Revisiting this idea in a new theoretical framework \\cite{add}, we proposed models that achieve interestingly large cross sections \\cite{us1}. We speculated that the GZK ``barrier'' could be then broken by neutrinos. The next question to answer is whether the shower events predicted look like events observed. Or are the characteristics so different from observed showers, which are generally compared to those of the simulations of proton, nucleus and gamma initiated showers, that the neutrino can be eliminated as a candidate for the super GZK showers? {\\it Our conclusion based on this study is no, they cannot be eliminated as candidates}.  For a range of values of the fundamental scale $M$ in the neighborhood of $2 -3$ TeV, there are neutrino energies 25\\% - 75\\% above that of the comparison proton model where the simulations match quite closely.  As noted in the introduction, the same analysis applies to a variety of cases - different interactions and different identities of primary particle.  Given that the same cross section input is not unique to the low scale gravity inspiration used here, this result is of quite general use.  Whether neutrinos follow the standard model extrapolations \\cite{q+r,Frich96}, are enhanced ``modestly'' by 3-5 orders of magnitude or ``extravagantly'' by more than 5 orders of magnitude, the search for UHE neutrino induced events at present and new facilities \\cite{auger,icecube} will be an exciting one.  {\\bf Note Added}: As we were completing this paper, a closely related work appeared \\cite{neu}.  The cross sections considered in their work are well below 10 mb, while we concentrate primarily on cross sections above this value.  Where cross sections are roughly the same, results and conclusions qualitatively agree.   {\\bf Acknowledgments:} DM thanks Seif Randjbar-Daemi and the High Energy Group at Abdus Salam I.C.T.P. and Jack Gunion and the High Energy Group at the Department of Physics, U.C. Davis for sabbatical leave hospitality while this work was being completed. AJ and PJ thank S. J. Sciutto for help in using the AIRES air shower simulator. We also thank Jaime Alvarez for useful discussions.  This work was supported in part by U.S. DOE Grant number DE-FG03-98ER41079, the {\\it Kansas Institute for Theoretical and Computational Science} and DST (India) grant No. DST/PHY/19990184."
    },
    "0011/astro-ph0011249_arXiv.txt": {
        "abstract": "The detection of Gamma Ray Burst GRB990705 on 1999, July 5.66765 UT, pointing to the Large Magellanic Clouds, suggested the search for a possible neutrino counterpart, both in coincidence with and slightly before (or after) the photon burst.  We exploited such a possibility by means of the LVD neutrino telescope (National Gran Sasso Laboratory, Italy), which has the capability to study low-energy cosmic neutrinos. No evidence for any neutrino signal, over a wide range of time durations, has been found, at the occurrence of GRB990705. Due to the lack of information about both the source distance and its emission spectrum, the results of the search are expressed in terms of upper limits, at the Earth, to the $\\bar{\\nu}_\\mathrm{e}$ flux $\\cdot$ cross-section, integrated over different time durations, $\\int \\int \\Phi_{\\bar \\nu_\\mathrm{e}}\\sigma dE dt$.  Moreover, assuming thermal $\\bar\\nu_\\mathrm{e}$ spectra at the source, upper limits to the $\\bar\\nu_\\mathrm{e}$ flux, integrated over time duration, for different spectral temperatures, are obtained. Based on these limits and on the expectations for $\\nu$ emission from collapsing astrophysical objects, the occurrence of a gravitational stellar collapse can be excluded up to a distance $r \\approx 50$ kpc, in the case of time coincidence with GRB990705, and $r \\approx 20$ kpc, for the 24 hours preceding it. ",
        "introduction": "Gamma Ray Burst GRB990705 was detected on 1999, July 5.66765 UT, by the BeppoSAX Gamma-Ray Burst Monitor, and localized by the BeppoSAX Wide Field Camera (Celidonio et al., 1999). It was promptly noted (Djorgovski et al., 1999) that its position, in projection, corresponded to the outskirts of the Large Magellanic Cloud (LMC), and it was suggested that, if the burst was indeed located in the LMC or its halo, a search for a neutrino signal, coincident with, or just prior to the GRB, would be quite interesting.  At the time of GRB990705, the LVD neutrino observatory, located in the Gran Sasso underground Laboratory, Italy, was regularly taking data, with active scintillator mass $M=573$ tons. The main purpose of the telescope is the search for neutrinos from gravitational stellar collapses in the Galaxy.  On July 19$^{\\mathrm{th}}$ 1999, the result of a preliminary analysis of the LVD data recorded during 48 hours around the time of GRB990705 was reported (Fulgione, 1999), and the absence of a neutrino signal, that would be expected from a gravitational stellar collapse in our Galaxy, was established (no additional results from other neutrino observatories were reported).  The search for low-energy neutrinos possibly associated to GRBs is indeed of interest, especially in view of the recent observational evidence linking (some) GRBs and supernovae (see, e.g., Galama et al., 1998, Bloom et al., 1999, Reichart, 1999). Many recent widely discussed models of the sources of GRBs involve the core collapse of massive stars (see, e.g., Woosley, 1993, Paczynski, 1998, Mac Fayden \\& Woosley, 1999, Khokhlov et al., 1999, Wheeler et al., 1999): in this scenario the neutrino emission could be associated to the cooling phase of the collapsed object, the time separation between the neutrino and gamma signals depending on the time necessary to transfer energy from the central engine, which emits thermal $\\nu$, to the outer region, emitting high energy photons.  It is clear that the possibility of detecting neutrinos correlated to GRBs depends on the distance of the associated source: even if it appears established that most of them lie at cosmological distances (Metzger et al., 1997), there is evidence, for at least one of GRBs, to be related to a supernovae event in the local universe (Tinney et al., 1998). In particular, from the study of the afterglow of GRB990705 (Masetti et al., 2000), although an extragalactic origin might be supported, the association with LMC cannot be ruled out.  Consequently, a more careful analysis of the LVD data in correspondence of GRB990705 has been performed, to search for weaker neutrino signals, not only in coincidence with, but also preceding\\footnote{ By analogy with SN explosions modelling, where few hours are required for the shock to reach the star envelope and give rise to the sudden increase of luminosity, a similar time gap can be assumed, between neutrinos and high-energy $\\gamma$-rays. } and even shortly following it.  The paper is planned as follows: in Sect.2 we briefly describe the LVD detector, and we explain the structure of the data. In Sect.3 we present the results of the analysis: a search for a $\\bar{\\nu}_\\mathrm{e}$ signal coincident in time with GRB990705 has been performed. Moreover, a time interval spanning from 24 hrs preceding the burst up to 10 minutes later, has been scanned, searching for any non-statistical fluctuation of the background. For sake of completeness, a wider interval, since 10 days before to 1 day after the event, has been investigated. We conclude in Sect.4, discussing the results in terms of upper limits to the $\\bar \\nu_\\mathrm{e}$ flux possibly associated to the GRB, under the hypothesis of thermal neutrino energy spectrum at the source, and comparing such limits with the expectations from existing models on $\\nu$ emission from collapsing objects. ",
        "conclusions": "%   The number of expected $\\bar\\nu_\\mathrm{e}$ interactions, $N_{\\mathrm{ev}}$, in a time interval $\\delta t$, due to a pulsed $\\bar\\nu_\\mathrm{e}$ emission, is defined as: $$N_{\\mathrm{ev}} = M  \\cdot N_\\mathrm{p} \\cdot \\epsilon \\int\\limits^{\\delta t}_{0} dt \\int\\limits^\\infty_{E_{\\mathrm{min}}} \\frac{d^2\\phi_{\\bar\\nu_\\mathrm{e}}}{dE_{\\bar\\nu_\\mathrm{e}} dt} \\sigma(E_{\\bar\\nu_\\mathrm{e}}) dE_{\\bar\\nu_\\mathrm{e}}$$ where $\\epsilon$ is the detection efficiency, $M$ [ton] is the active scintillator mass, $N_\\mathrm{p}$ is the number of free protons per scintillator ton, $\\sigma(E_{\\bar\\nu_\\mathrm{e}})$ is the neutrino interaction cross section (Vogel, 1984) and $\\frac{d^2 \\phi_{\\bar\\nu_\\mathrm{e}}}{dE_{\\bar\\nu_\\mathrm{e}} dt}$ is the differential neutrino flux at the Earth.  In the absence of any information on the source distance and its emission spectrum, we can express the results of the search in terms of upper limits to the flux $\\cdot$ cross-section, integrated over the time duration, at the Earth: $ \\int dt \\int \\frac{d^2\\phi}{dE dt} \\sigma dE$.  These limits, calculated at 90\\% c.l., are reported in Table 2, for various burst duration $\\delta t$, and they are expressed in number of interactions per target proton.   Any hypothesis on the $\\bar\\nu_\\mathrm{e}$ source spectrum leads to a limit to the time integrated $\\bar\\nu_\\mathrm{e}$ flux at the Earth. Assuming a thermal spectrum, constant during the emission interval $\\delta t$, i.e.: $$\\frac {d\\Phi_{{\\bar\\nu_\\mathrm{e}}}}{dE_{\\bar\\nu_\\mathrm{e}}} \\propto \\frac {(\\frac{E_{\\bar\\nu_\\mathrm{e}}}{T_{\\bar \\nu_\\mathrm{e}}})^2} {1+exp(- \\frac{E_{\\bar\\nu_\\mathrm{e}}}{T_{\\bar \\nu_\\mathrm{e}}})}$$ upper limits to the time integrated $\\bar\\nu_\\mathrm{e}$ flux are obtained, as a function of the neutrinosphere emission temperature $T_{\\bar \\nu_\\mathrm{e}}$ [MeV]. These results are shown in Fig.\\ref{fig:l2}, for burst duration $\\delta t \\leq 10$ s.  \\begin{figure} \\mbox{\\epsfig{file=ms10019.f3,height=8cm,width=8cm}}   % \\caption{Upper limits ($90 \\%$ c.l.) to the time integrated $\\bar\\nu_\\mathrm{e}$ flux, at the Earth, for thermal $\\bar \\nu_\\mathrm{e}$ spectra and $\\delta t \\leq 10$ s, compared with expectations for different source distances.} \\label{fig:l2} \\vspace{-0.5cm} \\end{figure}  Most theoretical models on the $\\bar\\nu_\\mathrm{e}$ emission from gravitational stellar collapses (Burrows, 1992) predict that the neutron star binding energy, $E_\\mathrm{b} \\approx 3 \\cdot 10^{53}$ erg, is emitted in neutrinos of every flavour (energy equipartition) with thermal energy spectra, during a time interval $\\delta t \\approx 10$ s. The corresponding $\\bar\\nu_\\mathrm{e}$ fluxes at the Earth, calculated, under the approximation of isotropical emission and pure Fermi-Dirac spectrum, for two different source distances: $50$ kpc (i.e., corresponding to the LMC\\footnote {One can compare these results with the neutrino flux observed from SN1987A, which was definitely located in the LMC. According to the combined analysis of the events detected by the KamiokandeII and IMB detectors (Jegerlehner et al., 1996), which yields a total emitted energy $E_\\mathrm{b} = 3.4 \\cdot 10^{53}$ erg and a $\\bar\\nu_\\mathrm{e}$ spectral temperature $T_{\\bar \\nu_\\mathrm{e}} = 3.6$ MeV, the resulting $\\bar\\nu_\\mathrm{e}$ flux at the Earth, integrated over time, is $\\Phi(\\bar\\nu_\\mathrm{e})\\cdot \\delta t \\sim 9. \\cdot 10^{9} \\mathrm {cm}^{-2}$ } ) and $20$ kpc (i.e., corresponding to the outskirts of our Galaxy), are reported in Fig.\\ref{fig:l2} and are compared with the results of the burst search.  The occurrence of a gravitational stellar collapse, with $\\bar \\nu_\\mathrm{e}$ emitted in the temperature range $T_{\\bar\\nu_\\mathrm{e}} > 2$ MeV, can then be excluded within a region of radius $r\\approx50$ kpc, in the case of time coincidence with the GRB990705 event, and $r\\approx20$ kpc, for the 24 hours preceding the GRB time\\footnote{ A possible effect of neutrino mixing on the signal from a gravitational stellar collapse would result in the merging of the energy spectra of neutrinos of different flavours. Because we are dealing with electron antineutrinos, which are characterized by a spectral temperature lower then the one of $\\bar\\nu_{\\mu}$ and $\\bar\\nu_{\\tau}$, neutrino oscillation effects would lead to a hardening of the $\\bar\\nu_\\mathrm{e}$ spectrum and, after all, to an increase of the $\\bar\\nu_\\mathrm{e}$ detection probability. Therefore, excluding oscillations into sterile neutrinos, the limits obtained in this work would remain valid even in the case of neutrino mixing.}."
    },
    "0011/astro-ph0011280_arXiv.txt": {
        "abstract": "We present a list of  34 neglected entries from star cluster catalogues located at relatively high galactic latitudes ($|b| >$ 15$^{\\circ}$) which appear to be candidate late stages of star cluster dynamical evolution.  Although underpopulated with respect to usual open clusters, they still present a high number density contrast as compared to the galactic field. This was verified by means  of (i) predicted model counts from different galactic subsystems in the same direction, and (ii) Guide Star Catalog equal solid angle counts for the object and surrounding fields. This suggests that the  objects are physical systems, possibly  star clusters in the process of disruption or  their fossil remains. The sample will be useful for followup studies in view of  verifying their physical nature. ",
        "introduction": "Star Clusters are known to dynamically evolve and stellar depletion effects eventually  lead to the cluster dissolution. Fundamental questions are: (i) where are the clusters in process of dissolution? (ii) if  fossils are left, can any remnant  be detected? The present study aims at showing that several candidates for these effects occur in star cluster catalogues themselves.  Several poorly populated  objects  at relatively high galactic latitudes (from the  open cluster perspective $|b| >$ 15$^{\\circ}$) are included in open cluster catalogues (Alter et al. 1970, Lyng\\aa  \\,1987). Other objects were reported as clusters  in  early studies (e.g. New General Catalogue and Index Catalogue), or in modern ones like the ESO (B) Atlas Survey Catalogue (Lauberts 1982).  Depletion of Main Sequence (MS) stars  has been detected or evidence of it has been found in luminosity functions and Colour-Magnitude Diagrams (CMDs) of some Palomar or Palomar-like  globular clusters such as E3 (McClure et al 1985), ESO452-SC11 (Bica et al. 1999) and NGC6717 (Palomar 9) (Ortolani et al. 1999). Such evolved dynamical stages of low mass globular clusters are still associated to relatively well  populated star clusters, but one may wonder what subsequent stages would look like, probably an underpopulated fossil core containing some double and multiple stars.  Among open clusters low  MS   depletion has been found e.g. in the intermediate age ($\\approx$ 3 Gyr)  cluster NGC3680 (Anthony-Twarog et al. 1991). The dissolution of open clusters has been studied by Wielen (1971). Updated data suggest that most open clusters dissolve in about 100 Myr and this will probably not change much as one includes fainter clusters  (Ahumada et al. 2000). Intermediate age open clusters are certainly survivors of initially massive clusters  (Friel 1995), and their  updated age histogram containing more than 100 entries (Dutra \\& Bica 2000) suggests a dissolution timescale of about 1 Gyr. Also, N-body simulations of star clusters in an external potential have shown typical dissolution times in the range 500-2500 Myr (Terlevich 1987, McMillan \\& Hut 1994, de la Fuente Marcos 1997, Portegies Zwart et al. 2000). Mass segregation is expected to occur in cluster cores during one relaxation time, according to N-body simulations (Terlevich 1987, Portegies Zwart et al. 2000). An important effect of mass segregation is the depletion of low mass MS stars by means of evaporation due to the tidal field of the Galaxy and encounters with binary stars. This would imply that clusters which are close to disruption have a core rich in compact and giant stars (Takahashi \\& Portegies Zwart 2000).  Recently, evidence of an open cluster remnant has been discussed by Bassino et al. (2000). They studied the relatively high latitude concentration of stars M73 (NGC6994) and derived an age of 2-3 Gyr from CMDs. They found a significant number density contrast with respect to the galactic field  CMD predicted by count models in the area.  Carraro  (2000) does not favour the object as an open cluster or as a remnant. At any rate, if M73 is a physical system the open cluster classification is certainly not adequate. Let us then suggest the acronym POCR - Possible Open Cluster Remnant. As an ongoing study of this neglected class of interesting objects and in view of future  CMDs to determine parameters such as reddening, age and distance, we present a list of  candidates. We  discuss their possible physical nature by checking whether  they present a significant number density contrast with respect to their fields. In Section 2 we present the sample. In Section 3  we analyze the significance of the excesses of  stars by means of (i) equal solid angle counts in the object and field areas, and (ii) galactic model counts. In Section 4 we discuss the results. Finally, in Section 5 we present the concluding remarks. ",
        "conclusions": "We presented a list of 34 neglected entries from star cluster catalogues located at relatively high galactic latitudes ($|b| >$ 15$^{\\circ}$)  which may be late stages of cluster dynamical evolution.  Although underpopulated with respect to usual open clusters, we showed that they still present a high number density contrast with respect to the galactic field, as verified by means  of (i) predicted model counts from different galactic subsystems in the same direction, and (ii) Guide Star Catalog equal solid angle counts for the object and surrounding fields. The sample will be useful for followup studies, aimed at  verifying their physical nature. Photometry and spectroscopy are required to determine fundamental parameters such as reddening, distance, age, radial velocities, membership and chemical abundances. Some of these objects might be clusters in the process of disruption or  their fossil remnants. The dynamical state  of the physical objects in the present sample may be inferred from comparisons of the cluster remnant and surrounding field luminosity functions, searching  for depletion effects.  An important population of possible open cluster remnants is likely to  exist. They may survive significant amounts of time as depopulated systems before dissolution. Simple arguments based on available numerical models and catalogued open clusters suggest that several hundreds can be expected. Systematic surveys to find new candidates and numerical models to explore in detail the evolved dynamical stages are encouraged for a better understanding of this so far overlooked object class."
    },
    "0011/astro-ph0011555_arXiv.txt": {
        "abstract": "We present an overview of the upcoming Microwave Anisotropy Probe (MAP) mission, with an emphasis on those aspects of the mission that simplify the data analysis.  The method used to make sky maps from the differential temperature data is reviewed and we present some of the noise properties expected from these maps.  An overview of the method we plan to use to mine the angular power spectrum from the mega-pixel sky maps closes the paper. ",
        "introduction": "In 1992 NASA's Cosmic Background Explorer (COBE) satellite made a full sky map of the cosmic microwave background (CMB) temperature with 7$^\\circ$ resolution, uncorrelated pixel noise, minimal systematic errors, and accurate calibration, from which CMB temperature anisotropy was first discovered \\cite{smoot92},  \\cite{bennett92}, \\cite{wright92}, \\cite{bennett96}.  The purpose of the MAP mission is to re-map the anisotropy over the full sky with more than 30 times the angular resolution ($\\sim 0.23^{\\circ}$ FWHM) and more than 35 times the sensitivity  ($\\sim 20$ $\\mu$K per 0.3$^{\\circ}$ pixel) of COBE, but with the same level of quality control as was possible with COBE. With this data, MAP will measure the physical interactions of the photon-baryon fluid (sound waves) in the early universe and thereby test models of structure formation, the geometry of the universe, and inflation.  In the years since COBE, a host of ground-based and balloon-borne experiments have detected and characterized fluctuations at smaller angular scales, most recently the experiments TOCO \\cite{toco}, BOOMERanG \\cite{boomldb} and MAXIMA \\cite{maxima}. However, because of their proximity to the Earth and its atmosphere, none of the ground or balloon-based experiments enjoy the extent of systematic error rejection or calibration accuracy that was possible with COBE. Moreover, many of these experiments have significantly correlated noise that places severe demands on the data analysis.  The COBE data still serves as a benchmark for the field and many aspects of the COBE mission have influenced the design of the MAP mission.  The need to minimize the level of systematic errors has been the major driver of the MAP design and has led to the following high level design features:  \\begin{itemize} \\item {\\bf A highly symmetric differential design}: MAP is a differential experiment based on pseudo-correlation microwave  radiometers that employ phase-matched HEMT amplifiers.  The instrument measures temperature differences between two  points $\\sim$ 141$^{\\circ}$ apart on the sky.  By measuring temperature differences, rather than absolute temperatures, many spurious signals will be common-mode and thus cancel upon differencing.  Also, by employing a pseudo-correlation design with a fast chopping frequency between the two sky inputs, 1/f noise that arises from the HEMTs can be chopped out. The resulting power spectrum of the radiometer noise is very nearly white (see \\S 3).  \\item {\\bf Multi-frequency}: There are five frequency bands from 22-90 GHz that will allow emission from the Galaxy and other non-cosmological sources to be modeled and removed based on their frequency dependence.  In the lowest frequency bands, MAP will probe the high frequency tail of radio emission from our Galaxy and provide valuable data on the enigmatic microwave foreground emission that correlates with thermal dust emission, but has a much different spectrum.  See \\cite{kogut} for a recent summary of evidence for this foreground.  \\item {\\bf Stability}: MAP will observe from a Lissajous orbit about the L2 Lagrange point 1.5 million km from Earth. The L2 point offers an exceptionally stable environment and an unobstructed view of deep space, with the Sun, Earth, and Moon always in shadow behind MAP's Sun shield.  MAP's large distance from Earth protects it from near-Earth emission and other disturbances.  While observing at L2, MAP's Sun shield and solar panels maintain a fixed angle with respect to the Sun  to provide exceptional thermal and power stability.  \\item {\\bf Low beam sidelobe levels}: The MAP optical system was designed with the foremost goal of providing adequate angular resolution along the line of sight while at the same time rejecting stray light from other directions.  For example, the largest instantaneous signal due to radiation from the galactic plane spilling into a sidelobe is expected to be less than 2 $\\mu$K at 90 GHz. \\end{itemize}  An overview of the MAP satellite is shown in Fig.~\\ref{map_view}.  The major visible features of MAP include the back-to-back telescope optics with 1.4 $\\times$ 1.6 m primary mirrors and 1 m secondary mirrors, the passive thermal radiators which cool the HEMT amplifiers to $<$100 K, the hexagonal structure housing the spacecraft service modules, and the large solar panel array/Sun shield which keeps the instrument in full shade.  MAP weighs a total of 830 kg and stands $\\sim$4 m tall.  It will be launched in 2001 aboard a Delta 7425-10 expendable launch vehicle from the NASA Kennedy Space Center Eastern Test Range.  \\begin{figure} \\centering \\includegraphics*[width=0.8\\hsize,keepaspectratio]{hinshawf1.eps} \\caption{\\small An overview of the MAP satellite.} \\label{map_view} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011139_arXiv.txt": {
        "abstract": "We investigate current problems in obtaining reliable ages for old stellar systems based on stellar population synthesis modelling of their integrated spectra.  In particular, we address the large ages derived for the globular cluster 47~Tuc, which is at odds with its Color-Magnitude-Diagram (CMD) age.  Using a new age indicator, H$\\gamma_{\\sigma<130}$, which is particularly effective at breaking the degeneracy between age and metallicity, we confirm the discrepancy between the spectroscopic age and the CMD age of 47~Tuc, in that the spectroscopic age is much older. Nebular emission appears unlikely to be a source for weakening the observed Balmer lines. We then explore a number of key parameters affecting the temperature of Turn-Off stars, which are the main contributors to the Balmer lines for old metal-rich stellar populations.  We find that $\\alpha$-enhanced isochrones with atomic diffusion included not only provides a good fit to the CMD of 47~Tuc, but also leads to a spectroscopic age in better agreement with the CMD age. ",
        "introduction": "An estimate of the mean luminosity-weighted stellar age of an early-type galaxy represents a major step in unveiling its true star formation history. However, to derive reliable information about stellar ages from the integrated light of unresolved galaxies one must deal with the age-metallicity degeneracy problem, which affects not only integrated colours but also absorption line-strengths (Worthey 1994). Recently, new age-dating techniques based on the Balmer lines (e.g., Jones \\& Worthey 1995; Vazdekis \\& Arimoto 1999, thereafter VA99) have shown great promise in untangling the age-metallicity degeneracy.  These techniques should be tested and calibrated on the metal-rich Galactic globular clusters (GCs) for which, unlike elliptical galaxies, independent age estimates are possible (Gibson et al. 1999, thereafter G99) by means of the Color-Magnitude-Diagram (CMD) of their resolved stellar population.  The application of the new age-dating techniques to very high signal-to-noise ratio (S/N) spectra of metal-rich Galactic GCs has revealed two major concerns: {\\it i)} the obtained ages are unreasonably large ($>$20~Gyr) (Jones 1999; Cohen, Blakeslee \\& Ryzhov 1998; G99; VA99) and {\\it ii)} a severe disagreement is found between the spectroscopic and CMD ages of 47~Tuc (G99). The CMD-derived ages may be sensitive to the dating method employed (e.g., Alonso et al. 1997) and to the input physics of the theoretical isochrones used as a reference (e.g., Salaris \\& Weiss 1998, thereafter SW98). However, a variety of recent CMD-based age determinations for 47~Tuc have consistently found its age to lie within 9-12.5 Gyr (SW98; Gratton et al.~1997; Carretta et al.~2000; Liu \\& Chaboyer~2000).  These CMD-derived ages are substantially younger than the $>20$~Gyr spectroscopically derived age obtained by G99 using the age indicator of Jones \\& Worthey (1995), and Worthey (1994) stellar population models. This discrepancy shows clearly that current stellar population synthesis models used for interpreting the integrated light of stellar systems may have severe zero points problems.  In \\S~2 we improve the spectroscopic age-dating technique which confirms a large spectroscopic age for 47~Tuc. In \\S~3 we discuss the possible origin of the problem, explore several theoretical parameters and suggest a possible solution. Finally, in \\S~4 we present our conclusions. ",
        "conclusions": "We have discussed the origin of the discrepancy between the spectroscopic (based on the effective temperature of TO stars) and CMD (based on the luminosity of TO stars and an assumed distance scale) age estimate for 47~Tuc as raised by G99. For this purpose we have defined a new age indicator, H$\\gamma_{\\sigma<130}$, particularly suitable for studying GCs and low velocity dispersion galaxies, which shows a superb power to break the age-metallicity degeneracy. H$\\gamma_{\\sigma<130}$ confirms the age discrepancy found by G99 for 47~Tuc. Emission fill-in of the Balmer lines appears to be an unlikely source of the weak H$\\gamma$ in 47~Tuc, since the ages derived from different Balmer lines give discordant results if the hypothetical emission fill-in is corrected for, and since a composite of four other metal-rich Galactic globular clusters shows the same weak H$\\gamma$ phenomenon. Thus the fact that other metal-rich GCs show very similar low Balmer values in comparison to the model predictions suggests a problem in the zero point of current stellar population models. It is worth noting that this zero point problem of the models with respect to the metal-rich GCs also works out for old elliptical galaxies.  We therefore analyzed the possible causes of the problem by studying a number of input parameters of the evolutionary computations, and comparing the observed value of H$\\gamma_{\\sigma<130}$ with that derived from the synthesized integrated spectra. Neither the initial He content nor the HB have significant effects on the Balmer indices synthesized for 47~Tuc. However the inclusion of $\\alpha$-elements enhancement and atomic diffusion in the evolutionary models provide spectroscopic ages which are much closer to the CMD derived ages. This occurrence constitutes a possible solution to the age-discrepancy between CMD and integrated spectrum ages of old metal-rich stellar populations.  It is important to study if the age discrepancy is present in metal poor GCs. For this purpose we need to expand the current stellar spectral libraries which feed the stellar population models (see V99), to extend their predictions to lower metallicities."
    },
    "0011/astro-ph0011413_arXiv.txt": {
        "abstract": "We present HST images at 622\\,nm and 300\\,nm of the jet in 3C\\,273 and determine the run of the optical spectral index at 0\\farcs2 along the jet.  We find no evidence for localized acceleration or loss sites, and support for a little-changing spectral shape throughout the jet. We consider this further evidence in favour of a distributed acceleration process. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011375_arXiv.txt": {
        "abstract": "We present an update of results from the search for microlensing towards the Large Magellanic Cloud (\\lmc) by \\eros\\ (Exp\\'erience de Recherche d'Objets Sombres). We have now monitored 25 million stars over three years.  Because of the small number of observed microlensing candidates (four), our results are best presented as upper limits on the amount of dark compact objects in the halo of our Galaxy.  We discuss critically the candidates and the possible location of the lenses, halo or \\lmc . We compare our results to those of the {\\sc macho} group.  Finally, we combine these new results with those from our search towards the Small Magellanic Cloud as well as earlier ones from the {\\sc eros1} phase of our survey.  The combined data is sensitive to compact objects in the broad mass range $10^{-7} - 10 \\, {\\rm M}_{\\odot}$. The derived upper limit on the abundance of stellar mass {\\sc macho}s rules out such objects as the dominant component of the Galactic halo if their mass is smaller than  $2 {\\rm M}_{\\odot}$. ",
        "introduction": "The search for gravitational microlensing in our Galaxy has been going on for a decade, following the proposal to use this effect as a probe of the dark matter content of the Galactic halo \\cite{pac86}.  The first microlensing candidates were reported in 1993, towards the \\lmc\\ \\cite{aub93,alc93} and the Galactic Centre \\cite{uda93} by the \\eros , \\macho\\ and \\ogle\\ collaborations.  Because they observed no microlensing candidate with a duration shorter than 10~days, the \\erou\\ and \\macho\\ groups were able to exclude the possibility that more than 10\\% of the Galactic dark matter resides in planet-sized objects \\cite{aub95,alc96,ren97,ren98,alc98}.  However a few events were detected with longer time\\-scales. In their two-year analysis \\cite{alc97a}, the \\macho\\ group estimated, from 6-8 candidate events towards the \\lmc , an optical depth of order half that required to account for the dynamical mass of a ``standard'' spherical dark halo\\footnote{ $4 \\times 10^{11}\\:{\\rm M}_\\odot$ within 50~kpc}; the typical Einstein radius crossing time of the events, $t_E$, implied an average mass of about 0.5~M$_\\odot$ for halo lenses \\cite{alc97a}. Based on two candidates, \\erou\\ set an upper limit on the halo mass fraction in objects of similar masses \\cite{ans96,ren97}, that is below that required to explain the rotation curve of our Galaxy\\footnote{ Assuming the original two \\erou\\ candidates are microlensing events, they would correspond to an optical depth six times lower than that expected from a halo fully comprised of \\macho s.}.  The second phase of the \\eros\\ programme was started in 1996, with a ten-fold increase over \\erou\\ in the number of monitored stars in the Magellanic Clouds. The analysis of the first two years of data towards the Small Magellanic Cloud (\\smc ) allowed the observation of one microlensing event \\cite{pal98} also detected by \\cite{alc97b}. This single event, out of 5.3 million monitored stars, allowed \\erod\\ to further constrain the halo composition, excluding in particular that more than 50~\\% of the standard dark halo is made up of $0.01 - 0.5 \\:{\\rm M}_\\odot$ objects \\cite{afo99}. In contrast, an optical detection of a halo white dwarf population was reported \\cite{iba99}, compatible with a galactic halo full of white dwarfs.  Very recently, the \\macho\\ group presented an analysis of 5.7 year light curves of 10.7~million stars in the \\lmc\\ \\cite{alc00} with an improved determination of their detection efficiency and a better rejection of background supernova explosions behind the \\lmc . They now favour a galactic halo \\macho\\ component of 20\\% in the form of 0.4~M$_\\odot$ objects. Within a few days, the detection of a halo white dwarf population at the level of a 10\\% component was also reported by \\cite{iba00}. Simultaneously, the \\erod\\ group presented its results from a two-year survey of 17.5 million stars in the \\lmc\\ \\cite{las00a}. One \\erou\\ microlensing candidate, {\\sc eros1-lmc-2}, was seen to vary again, 8 years after its first brightening, and was thus eliminated from the list of microlensing candidates. Two new candidates were identified ({\\sc eros2-lmc-3} and~4). Because this is much lower than expected if \\macho s are a substantial component of the galactic halo, and because these two new candidates do not show excellent agreement with simple microlensing light curves, \\eros\\ chose to combine these results with those from previous \\eros\\ analyses, and to quote an upper limit on the fraction of the galactic halo in the form of \\macho s.  In this article, we describe an update on the \\erod\\ \\lmc\\ data, an analysis of the three-year light curves from 25.5 million stars.  While the sensitivity is improved, the main conclusions are unchanged compared to \\cite{las00a}. One of the two-year candidates was seen to vary in the third season and was thus rejected.  Three new candidates have been detected. We combine these \\erod\\ \\lmc\\ results with those of previous independent \\eros\\ analyses, and derive the strongest limit obtained thus far on the amount of stellar mass objects in the Galactic halo. ",
        "conclusions": "After nine years of monitoring the Magellanic Clouds, \\eros\\ has a meager crop of five microlensing candidates towards the \\lmc\\ and one towards the \\smc , whereas about 30 events are expected towards the \\lmc\\ for a spherical halo fully comprised of $0.4 \\:{\\rm M}_\\odot$ objects. Moreover, some of the candidates cannot be considered excellent. These candidates were obtained from four different data sets analysed by independent, cross-validated programs. So, the small number of observed events is unlikely to be due to bad (and overestimated) detection efficiencies.  This allows us to put strong constraints on the fraction of the halo made of objects in the range [$10^{-7}\\:{\\rm M}_\\odot$, $10\\:{\\rm M}_{\\odot}$], excluding in particular at the 95\\% C.L. that more than 40\\% of the standard halo be made of objects with up to $1 \\:{\\rm M}_\\odot$. The preferred value quoted in \\cite{alc00}, $f = 0.2$ and $0.4\\:{\\rm M}_\\odot$, is consistent with our limit as can be seen in Fig.~\\ref{excl}. (The upper part -~about 25\\%~- of the domain allowed by \\cite{alc00} is excluded by the limit we report here.)  There are several differences which should be kept in mind while comparing the two experiments. First, \\eros\\ uses less crowded fields than \\macho\\ with the result that blending is relatively unimportant for \\eros . (Were \\eros\\ results to be corrected for blending, the detection efficiency would increase slightly and the reported limit would be stronger.) Second, \\eros\\ covers a larger solid angle (64~deg$^2$ in the \\lmc\\ and 10~deg$^2$ in the \\smc ) than \\macho , which monitors primarily 15~deg$^2$  in the central part of the \\lmc . The \\eros\\ rate should thus be less contaminated by self-lensing, i.e. microlensing of \\lmc\\ stars by dimmer \\lmc\\ objects, which should be more common in the central regions. The importance of self-lensing was first stressed in \\cite{wu94,sa94}. Third, the \\macho\\ data have a more frequent time sampling.  The results from \\eros\\ and \\macho\\ are apparently consistent, but the way they are interpreted is different.  \\macho\\ reports a signal and considers the contamination of its sample as low or null.  \\erod\\ quotes an upper limit and does not claim its sample to be background-free.  The position of the lenses along the line of sight, halo or Magellanic Clouds, is also an issue.  \\macho\\ has compared the spatial distribution of its candidates across the face of the \\lmc\\ and observes a better agreement with the halo hypothesis than with a specific model of the \\lmc .  On the other hand, because the \\eros\\ stars are spread over a wider field, the fact that the \\eros\\ sample corresponds to a lower central value of the event rate (about twice lower than that of \\macho ) is compatible with an interpretation where a noteable fraction of the events are due to self-lensing.  The small number of \\eros\\ candidates precludes at present any definitive conclusion on that topic.  It seems likely that the single most important input to the question of the position of the lenses will come from the comparison of the microlens candidates samples towards the \\smc\\ with those towards the \\lmc .  Because the two lines of sight are rather close (about 20~degrees apart), the timescale distributions of microlensing candidates towards the two Clouds should be nearly identical if lenses belong to the galactic halo.  Also, the event rates should be comparable, although the ratio is more halo model dependent. At present, \\eros\\ has analysed two seasons of \\smc\\ data \\cite{afo99} and \\macho\\ has not yet presented its detection efficiency towards the \\smc .  From the published \\eros\\ efficiencies, and assuming that the \\macho\\ efficiencies towards the \\smc\\ are similar to those towards the \\lmc , it can be expected that the completed experiments will have gathered between five and ten microlenses towards the \\smc .  This should allow a significant comparison of the timescales (see also the discussion in \\cite{gra00}).  Finally, let us mention that, given the scarcity of our candidates and the possibility that some observed microlenses actually lie in the Magellanic Clouds, \\eros\\ is not willing at present to quote a non zero {\\it lower} limit on the fraction of the Galactic halo comprised of dark compact objects with masses up to a few solar masses.  \\bigskip {\\bf Acknowledgements} We are grateful to D. Lacroix and the staff at the Observatoire de Haute Provence and to A. Baranne for their help with the MARLY telescope. The support by the technical staff at ESO, La Silla, is essential to our project. We thank J.F. Lecointe for assistance with the online computing. We thank all non-\\eros\\ members who participated in data taking."
    },
    "0011/astro-ph0011143_arXiv.txt": {
        "abstract": "\\emph{Eddington} is a space mission for extrasolar planet finding and for asteroseismic observations. It has been selected by ESA as an F2/F3 reserve mission with a potential implementation in 2008-13. Here we describe \\emph{Eddington}'s capabilities to detect extrasolar planets, with an emphasis on the detection of habitable planets. Simulations covering the instrumental capabilities of \\emph{Eddington} and the stellar distributions in potential target fields lead to predictions of about 10,000 planets of all sizes and temperatures, and a few tens of terrestrial planets that are potentially habitable. Implications of \\emph{Eddington} for future larger scale missions are briefly discussed. ",
        "introduction": "mission} The \\emph{Eddington} mission is a space telescope designed for two primary goals: asteroseismic studies and extrasolar planet finding. Both goals will be achieved through the acquisition of high-precision wide-field photometry, with a temporal stability that is possible only from space. The basic design is an f/3 triple-reflecting telescope with a 1.2m diameter aperture, 0.6m$^2$ effective area, and a 3$\\deg$ diameter field of view imaged by a 20-CCD mosaic camera. The telescope will be launched into an orbit around the L2 Earth-Sun libration point, from which the entire sky is accessible for observations. The first 2 years of the mission will be dedicated primarily to asteroseismic studies, with pointings lasting up to two months to a variety of targets. The mission's second part will be dedicated to planet finding by the transit method, where the telescope will survey a single stellar field for 3 years. In total, about 500,000 stars will be surveyed by \\emph{Eddington} as potential hosts for planetary systems.  In October 2000, \\emph{Eddington} was selected by ESA as a reserve mission for the F2/F3 launch window (2009-2013). More details about the mission can be found in the \\emph{Eddington} Assessment study (Favata et al., 2000) and in the \\emph{Eddington} web-site (http://astro.esa.int/SA-general/Projects/Eddington/). Here we will give a short overview about its planet detection capabilities. ",
        "conclusions": ""
    },
    "0011/astro-ph0011469_arXiv.txt": {
        "abstract": "We report the first results from the BOOMERanG experiment, which mapped at 90, 150, 240 and 410 GHz a wide ($3\\%$) region of the microwave sky with minimal local contamination. From the data of the best 150 GHz detector we find evidence for a well defined peak in the power spectrum of temperature fluctuations of the Cosmic Microwave Background, localized at $\\ell = 197 \\pm 6$, with an amplitude of $(68 \\pm 8) \\mu K_{CMB}$. The location, width and amplitude of the peak is suggestive of acoustic oscillations in the primeval plasma. In the framework of inflationary adiabatic cosmological models the measured spectrum allows a Bayesian estimate of the curvature of the Universe and of other cosmological parameters. With reasonable priors we find $\\Omega = (1.07 \\pm 0.06)$ and $n_s = (1.00 \\pm 0.08)$ (68$\\%$C.L.) in excellent agreement with the expectations from the simplest inflationary theories. We also discuss the limits on the density of baryons, of cold dark matter and on the cosmological constant. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011233_arXiv.txt": {
        "abstract": "Chiral cosmic strings are naturally produced at the end of D-term inflation and they may have interesting cosmological consequences. As was first proved by Carter and Peter, the equations of motion for chiral cosmic strings in Minkowski space are integrable (just as for Nambu-Goto strings).  Their solutions are labeled by a function $k(\\sigma - t)$ where $t$ is time and $\\sigma$ is the invariant length along the string, and the constraints on $k$, which determines the charge on the string, are that $0 \\leq k^2 \\leq 1$. We review the origin of this parameter and also discuss some general properties of such strings which can be deduced from the equations of motion.  The metric around infinite chiral strings is then constructed in the weak field limit, and studied as a function of $k$.  We also consider the angular momentum of circular chiral loops, and extend previous work to consider the evolution and self-intersection properties of a more general family of chiral cosmic string loops for which $k^2(\\sigma-t)$ is not constant. ",
        "introduction": "In the last few years the scenario of structure formation from cosmic strings has become increasingly tenuous, since its predictions differ significantly from the new high accuracy measurements of the temperature fluctuations in the cosmic microwave background radiation.  Most studies of such observational consequences of strings have focused on structureless Nambu-Goto (NG) strings \\dc{James1,CHM,avelino,Levon,CMS} and global strings \\dc{Neil,Ruth}, and in each case the recent predictions are based on numerical simulations of the evolution of the string network postulated to form at the GUT phase transition. One should recall though that there are some unresolved and potentially important uncertainties in the simulations --- it is very difficult, for example, to resolve the very disparate scales which characterize the the network, as well as to deal with gravitational backreaction effect --- and hence a combination of numerical work with analytical modeling \\dc{James1,Levon,CMS} has also been used to make predictions from NG strings.   Our focus here is not on NG strings but rather on {\\em chiral cosmic strings}. These strings are a type of current carrying string \\dc{Witten} for which the world-sheet current $j^{i}$ is null; $$ j^{i}j_i = j^2 = 0. $$ (Here $i=(0,1)$ and the 2D world sheet metric $\\gamma_{ij}$ defined below raises and lowers indices.) One motivation for studying such chiral strings comes from the well known SUSY D-term inflation model.  In this model, strings are produced at the end of inflation \\dc{Rachel} so that both mechanisms contribute to producing density fluctuations.  However, the strings produced are chiral cosmic strings and not NG strings \\dc{DDT1}.  Hence in order to make predictions for the $C_l$'s from this `strings plus inflation' model, the evolution and cosmological consequences of chiral cosmic string networks must be understood. (There may exist models in which the strings formed at the end of inflation are NG ones, however this is not true of D-term inflation. In the case of `inflation plus NG strings', predictions may be found in \\dc{CHM2}.)     There are a number of differences between the properties of chiral cosmic strings and NG strings.  One such regards the evolution of the strings themselves: the null current on chiral strings can, as in the case of other current-carrying strings, lead to the formation of non self-intersecting stable loops called vortons\\footnote{As will be come clearer later, by a vorton we mean a stable loop of arbitrary shape that never self-intersects. This definition is different from that of Martins and Shellard \\dc{MS} who also require that these loops move non-relativistically, suggesting that otherwise the charge on the loops could be `thrown off'.  We are not able to comment on such a mechanism, however see \\dc{DDDD} for a discussion of the scattering of zero-modes from chiral strings.}. This is potentially catastrophic as the energy density in the chiral string network could quickly dominate the energy density in the universe if stable vortons are present.  It is therefore important to see if vortons are produced, and in section \\dr{s:kvar} we study the self-intersection properties of a family of chiral cosmic string loops.  Another difference between NG and chiral strings is that these line-like sources of energy generate different metrics about them (section \\dr{s:metric})\\footnote{I am aware that this comment disagrees with one I made in \\dc{Verb}!  I would like to thank Patrick Peter and Tanmay Vachaspati for pointing out an error in my previous determination of the metric.}.  One might therefore expect them to produce different perturbations in the matter and radiation through which they pass.   Recently a number of steps have been made which allow for a quantitative study of chiral cosmic string dynamics. First, a well defined unique 2D effective action exists for these strings \\dc{MS,CP}.  From this action it was shown, with suitable gauge choices, that the equations of motion are integrable in Minkowski space \\dc{CP} (see also \\dc{US,BP} for different presentations of the same result).  They are \\be \\frac{\\paa^2 \\vecx}{\\paa t^2} - \\frac{\\paa^2 \\vecx}{\\paa \\sigma^2} = 0 \\qquad \\Longrightarrow \\qquad \\vecx(t,\\si) = \\frac{1}{2}[\\veca(t+\\si) + \\vecb(t-\\si)], \\dle{eqnlecon} \\ee where $t$ is background time, and $\\sigma$ measures the invariant length or energy along the string as in the NG case \\dc{US}. The constraints are \\ba \\spr{\\veca}^2 \\= 1, \\dla{c1} \\\\ \\spr{\\vecb}^2  \\!\\!\\!& \\le &\\!\\!\\! 1, \\dla{c2} \\end{eqnarray} where for instance $\\spr\\veca(q)\\equiv d\\veca(q)/dq.$  If one defines \\be k^2:= \\spr{\\vecb}^2 \\dle{kddef} \\ee so that $k=k(t-\\sigma)$, then it can be shown that $k^2$ determines the conserved charge on the string (see also below).  Furthermore, if $k = $ constant $=1$ then this charge vanishes as required, since $\\spr{\\vecb}=1$ is just the Nambu-Goto limit. In reference \\dc{US}, the self-intersection properties of chiral cosmic string loops were also studied in the special case of $k =$ constant. In particular the strings were shown never to self-intersect for $k=0$: this case corresponds to maximal charge on the strings and to vorton solutions.   Here that work is extended, though we still consider Minkowski space (with metric $\\eta_{\\mu \\nu} = (+,-,-,-)$) throughout. First, for completeness, we indicate in section \\dr{s:2} how the equations of motion (\\dr{eqnlecon})-(\\dr{c2}) are obtained from the chiral action and how the charge mentioned above is defined. This necessarily follows parts of reference \\dc{US} rather closely, though a small error in that paper is corrected.  We also compare the chiral charge with the charges used for more general current carrying strings.  In section \\dr{s:angmmt} we summarize some properties of chiral cosmic strings which result from the equations of motion.  The metric around infinite chiral strings is then studied as a function of $k$ and we comment in possible consequences it may have for structure formation and CMB anisotropies from chiral cosmic strings. In section \\dr{ss:ang}, the effect of angular momentum on the motion of circular loops is considered by looking at the effective potential introduced in \\dc{CPG}. In section \\dr{s:kvar} we investigate the self-intersection properties of loops with non-constant $k$. Finally conclusions are given in section \\dr{s:conc}. ",
        "conclusions": "\\dle{s:conc}  In this paper we have attempted study and clarify a number of points regarding the evolution and gravitational properties of chiral cosmic strings.  As was summarized in section \\dr{s:2}, the crucial difference between the equations of motion for NG and chiral cosmic strings is the constraint on the vector $\\spr{\\vecb}$: for NG strings $\\spr{\\vecb}^2(\\eta) = 1$ $\\forall \\eta$, whereas for chiral strings $\\spr{\\vecb}^2(\\eta) (=k^2(\\eta)) \\leq 1$.  Equation (\\dr{Cr}) shows that $k^{2}(\\eta)$ determines the charge on the chiral string.  We saw in section \\dr{s:gen} that chiral strings with $k=0$ ($\\forall \\eta$) move along themselves and never self-intersect.  If the string forms a loop, the energy of this arbitrary shaped vorton is equipartitioned between tension and angular momentum.  The charge on the vortons is given by $C = m L_{{\\rm phys}}$ where $ L_{{\\rm phys}}$ is the constant physical length of the vorton.  Infinite straight chiral strings were studied in section \\dr{s:metric}.  We saw that the energy momentum tensor contains non-diagonal terms $T^{tz} \\neq 0$.  These represent the momentum along the string. Furthermore, $T^{tt} \\neq T^{zz}$ (if $k \\neq 1$) which is reminiscent of the situation which occurs with wiggly NG strings.  As a consequence of the form of $T^{\\mu \\nu}$, the weak-field metric around the string was shown to contain a $dt \\, dz$ term which means that photons (and relativistic particles) moving near the string are dragged in the direction of the string.  We also observed that there is a $k$-dependent deficit angle as well as a $k$-dependent Newtonian potential.  Regarding the evolution of a chiral cosmic string network (which could formed at the end of D-term inflation), it is important to understand whether or not the loops can self-intersect and then decay.  If they cannot decay, this would lead to a cosmological catastrophe as they would dominate the energy density of the universe.  In section \\dr{ss:ang} we studied the effective potential for the motion of a non self-intersecting circular loop for which $0 \\leq k < 1$.  In section \\dr{s:kvar} we considered loops with non-constant $k$: the physical reason for which one might expect $k$ not to be constant is that charge will build up as a result of self-intersections, and also fluctuate during the phase transition which forms the string. Analysis of specific form of $k(\\eta)$ (given via (\\dr{bgen})) showed that self-intersection is possible for these loops.  The ensuing numerical analysis showed that the self-intersection probability depends on the form of $k(\\eta)$ and is not uniquely determined by the charge $C$ of the loop. This unfortunately suggests that even if one were able to estimate $C$ for the strings in a chiral cosmic string network, this would not be sufficient to determine the self-intersection properties of the loops. As a further problem it still remains to understand the fate of the daughter loops.  A number of interesting questions remain to be studied.  Regarding the metric (section \\dr{s:metric}), it would be interesting to go beyond the weak-field approximation and also to study carefully the potential cosmological consequences of the $dt \\, dz$ term \\dc{StVa}.  This cross-term is the main difference between the metric for NG and chiral strings.  Concerning the evolution of a network of chiral cosmic strings it is clear that if the network is formed with $k(\\eta) = 0$ $\\forall \\eta$ and for all strings, then this leads to a cosmological catastrophe: this is the only case in which the answer for $P_{int}$ is unique and zero! --- the strings cannot self-intersect and are frozen.  Similar problems occur if this state is reached at anytime during the evolution of the network.  This vorton problem was studied in \\dc{CD} where it was noted that the quantum number $C$ should be larger for chiral strings than for strings with time- or space-like currents. However, work still needs to be done to see if $C$ is maximal or not \\dc{DP}.  If it is not maximal (i.e.\\ $k \\neq 0$ $\\forall \\eta$) it still remains to understand the ultimate fate of the daughter loops, and hence that of the network itself."
    },
    "0011/astro-ph0011011_arXiv.txt": {
        "abstract": " ",
        "introduction": "The black hole candidate GRS~1915+105 is known for its extremely complex X-ray behaviour, which is not observed in any other X-ray source. Belloni et al. (2000) succesfully described this behaviour as transitions between three basic spectral states: the soft states A and B, and the hard state C. GRS 1915+105 was discovered in 1994 to be the first galactic source to show superluminal motions of radio-emitting ejecta (Mirabel \\& Rodriguez 1994). Besides these big radio flares, the source also shows smaller radio oscillation events with periods around 20-40 minutes (Pooley \\& Fender 1997). ",
        "conclusions": ""
    },
    "0011/astro-ph0011227_arXiv.txt": {
        "abstract": "We present a spectrophotometric study based on VLT/FORS\\,I observations of one of the most metal-deficient blue compact dwarf (BCD) galaxies known, \\tol\\ ($Z$$\\sim$$Z$$_{\\odot}$/25). The data show that roughly half of the total luminosity of the BCD originates from a bright and compact starburst region located at the northeastern tip of a faint dwarf galaxy with cometary appearance. The starburst has ignited less than 4 Myr ago and its emission is powered by several thousands O7V stars and $\\sim$ 170 late-type nitrogen Wolf-Rayet stars located within a compact region with $\\la$500 pc in diameter. For the first time in a BCD, a relatively strong [Fe\\,V] $\\lambda$4227 emission line is seen which together with intense He II $\\lambda$4686 emission indicates the presence of a very hard radiation field in \\tol. We argue that this extraordinarily hard radiation originates from both Wolf--Rayet stars and radiative shocks in the starburst region. The structural properties of the low-surface-brightness (LSB) component underlying the starburst have been investigated by means of surface photometry down to 28\\ $B$ \\sbb. We find that, for a surface brightness level fainter than $\\sim 24.5$\\ $B$\\ \\sbb, an exponential fitting law provides an adequate approximation to its radial intensity distribution. The broad-band colors in the outskirts of the LSB component of \\tol\\ are nearly constant and are consistent with an age below one Gyr. This conclusion is supported by the comparison of the observed spectral energy distribution (SED) of the LSB host with theoretical SEDs. ",
        "introduction": "} \\subsection{Photometric data} Images of \\tol\\ in the broad-band filters Bessell $U$, $B$, $R$ were obtained with the {\\sf Fo}cal {\\sf R}educer and low-dispersion {\\sf S}pectrograph (FORS\\,1; see Moehler et al. 1995) attached to the VLT\\ UT1. The exposures were acquired under photometric conditions on May, 17th 1999 during a 6 night run allocated to guaranteed time observations (GTO) at an airmass ranging between 1.48 and 1.64. The seeing was between 0\\farcs 7 and 0\\farcs 9. FORS was operating in the standard imaging mode yielding a final focal ratio of 3.13 and an instrumental scale of 0\\farcs 2 pix$^{-1}$.  Photometric zero-points and color-dependent calibration terms were derived from exposures of the standard-star field Mark\\ A (Landolt 1992) taken each night during the GTO run at an airmass$\\approx$1. Airmass-dependent calibration terms were obtained using standard extinction curves and found to agree to a level better than 10\\% with those derived during the commissioning phase\\,II of FORS. The photometric precision is estimated to be $\\sim 0.1$ mag. Reduction has been accomplished in the standard way using the ESO-MIDAS software package. \\subsection{Spectroscopic data} Spectroscopic data for Tol 1214--277 were taken with FORS\\,1 on May, 12th 1999 at an airmass 1.7 and with a seeing between 0\\farcs7 -- 0\\farcs9 FWHM. A 1\\arcsec\\ $\\times$ 180\\arcsec\\ slit was used in conjunction with a grism GRIS$\\_$300V and a GG375 second-order blocking filter. This yields a spatial resolution along the slit of 0\\farcs2 pixel$^{-1}$, a scale perpendicular to the slit of 3 \\AA\\ pixel$^{-1}$, a spectral coverage of 3600 -- 7500 \\AA, and a spectral resolution of $\\sim$ 10 \\AA\\ (FWHM). The total exposure time of 1650 seconds has allowed to reach a signal-to-noise ratio S/N $\\ga$ 50 in the continuum of the bright central part of the BCD and was broken up into two subexposures, 990 and 660 seconds, to allow for an efficient cosmic-ray removal. The slit was oriented in the position angle P.A. = 39$^{\\circ}$ to enable a simultaneous study of the starburst knot and the faint underlying host along its major axis. A spectrum of a He-Ne-Ar comparison lamp was obtained for wavelength calibration.  Since no spectrophotometric standard star was observed, calibration was done using flux-calibrated spectra of \\tol\\ obtained previously with the 2.1m KPNO telescope\\footnote{Kitt Peak National Observatory (KPNO) is operated by the Association of Universities for Research in Astronomy (AURA), Inc. under cooperative agreement with the National Science Foundation.}. These were taken on April 2, 1998 at an airmass 2.0. The total exposure time was 3600 seconds, split into 3 subexposures, 1200 seconds each. A 2\\arcsec\\ $\\times$ 180\\arcsec\\ slit was used along with grating No. 9 and a GG375 second-order blocking filter. This yields a spatial resolution along the slit of 0\\farcs69 pixel$^{-1}$, a scale perpendicular to the slit of 2.7 \\AA\\ pixel$^{-1}$, a spectral range 3600 -- 7500 \\AA, and a spectral resolution of $\\sim$ 7 \\AA\\ (FWHM). The spectrophotometric standard star Feige 34 was observed for flux calibration. Data reduction of the 2.1m telescope and VLT spectroscopic observations was carried out using the IRAF software package. This included bias subtraction, cosmic-ray removal, flat-field correction, wavelength calibration and night-sky background subtraction. ",
        "conclusions": "The main conclusions drawn from our imaging and spectroscopic analysis of deep VLT data of the extremely metal-deficient ($Z$$\\sim$$Z$$_{\\odot}$/25) and nearby ($D$ = 103.9 Mpc) BCD \\tol\\ may be summarized as follows:  1. \\tol\\ undergoes a vigorous burst of star formation having ignited less than 4 Myr ago. The starburst takes place within a bright ($M_B\\!\\sim\\!-16$ mag) compact ($\\la$ 500 pc in diameter) region,  giving rise to extended and abundant ionized gas emission with an H$\\beta$ equivalent width of $\\sim$ 320 \\AA. The starburst is powered by several thousands of O7V stars and 170 late-type nitrogen Wolf--Rayet stars.  2. In this very metal-deficient BCD we discover for the first time the high ionization line [Fe V] $\\lambda$4227. Moreover, we detect extraordinarily strong He II $\\lambda$4686 emission with an intensity as high as 5\\% of that of the H$\\beta$ emission line. This implies the presence of a very hard radiation field in \\tol. The intensity ratio $I$([Fe V] $\\lambda$4227) / $I$(He II $\\lambda$4686) in \\tol\\ compares well with that in another extremely metal-poor BCD with [Fe V] $\\lambda$4227 emission, SBS 0335--052, being in both cases larger by more than one order of magnitude than the ratio observed in high-excitation planetary nebulae. While the relative number of WR stars of $N$(WR) / $N$(O + WR) = 0.023 in \\tol\\ is compatible with theoretical predictions, the intensity of the He II $\\lambda$4686 emission line exceeds several times the predictions of standard H II photoionization models, even when the hard radiation component of Wolf-Rayet stars is taken into account. Therefore, we argue that the hard ionizing radiation field in \\tol\\ is produced from the combined effect of massive stars and SN-driven shocks.  3. Star-forming activity in \\tol\\ is confined to the northeastern tip of a cometary dwarf galaxy with an absolute $B$ magnitude $\\ga$ --16 mag and an isophotal size of 7.6$\\times$4.8 kpc at 28 $B$ \\sbb. An exponential fitting law provides a reasonable approximation to the intensity distribution of the stellar LSB host in its outskirts, for $\\mu_B\\ga 24.5$ \\sbb. It fails, however, to properly describe the observed brightness distribution at intermediate and high intensity levels. These are better fitted by an exponential distribution which flattens at small radii, similar to the V-type profiles described by Binggeli \\& Cameron (1991).  4. The radially averaged ($U-B$) and ($B-R$) colors of the LSB host of \\tol\\ are consistent with those for either an instantaneous burst with log ($t$/yr) $\\la$ 8.1, or a continuous star formation between log ($t$/yr) $\\la$ 8.7 and log ($t$/yr) $=$ 7.3. We however cannot definitely exclude the presence of a small fraction of old (age $>$ 1 Gyr) stars due to their intrinsic faintness."
    },
    "0011/astro-ph0011157_arXiv.txt": {
        "abstract": "To determine the equation of state of the universe, we propose to use a new independent variable  $R\\equiv (H_0/c)(d_L(z)/(1+z))$, where $H_0$ and $d_L(z)$ are the present Hubble parameter and the luminosity distance, respectively. For the flat universe suggested from the observation of the anisotropy of cosmic microwave background, the density and the pressure are expressed as $\\rho/\\rho_0=4(df/dR)^2/f^6$ and $p/\\rho_0=-4/3(d^2f/dR^2)/f^5$ where $\\rho_0$ is the present density and $f(R)=1/\\sqrt{1+z(R)}$. In $(R, f)$ plane the sign as well as the strength of the pressure is in proportion to the curvature of the curve $f(R)$. We propose to adopt a Pade-like expression of $f(R)=1/\\sqrt{u}$ with $u\\equiv 1+\\sum\\limits_{n=1}^{N}u_nR^n$. For flat $\\Lambda$ model the expansion up to $N=7$ has at most an error $< 0.2\\%$ for $z < 1.7$ and any value of $\\Lambda$. We also propose a general method to determine the equation of state of the universe which has $N-1$ free parameters. If the number of parameters are smaller than $N-1$, there is a consistency check of the equation of state so that we may confirm or refute each model. ",
        "introduction": "Recent measurements of the luminosity distance $d_L(z)$ using Type Ia supernovae (\\cite{schmidt98,ries98,perl99}) suggest that accurate $d_L(z)$ may be obtained in the near future. Especially SNAP (\\cite{snap}) will give us the luminosity distance of $\\sim$2000 Type Ia supernovae with an accuracy of a few \\% up to $z\\sim1.7$ every year. On the other hand from the observation of the first Doppler peak of the anisotropy of CMB,  it is now suggested that the universe is flat (\\cite{ber00,lange00}), which may be proved  in future by MAP and Planck. If the flat universe is the case, the density $\\rho(z)$ and the pressure $p(z)$ can be determined only from  $d_L(z)$  in principle (\\cite{naka99}) so that the equation of the state of the universe is uniquely determined. If not, the determination of the present curvature of the universe  and the determination of $\\rho(z)$ and  $p(z)$ will be coupled in general (\\cite{naka99}).  Now let us assume that the universe is flat. Even in this case at least two problems exist: 1) How to express the continuous function $d_L(z)$ which is  accurate enough from $z=0$ to $z\\sim 1.7$ using several free parameters; 2) From $d_L(z)$ how to obtain accurate  $\\rho(z)$ and  $p(z)$, that is, the equation of state of the universe(\\cite{star98,huter99,naka99,saini99,chiba00}).  In this paper, we  propose to use a new independent variable  $R\\equiv (H_0/c)(d_L(z)/(1+z))$ instead of $z$, where $H_0$ is the present Hubble parameter. We show that  $\\rho(R)/\\rho_0=4(df/dR)^2/f^6$ and $p(R)/\\rho_0=-4/3(d^2f/dR^2)/f^5$ where $\\rho_0$ is the present density and $f(R)=1/\\sqrt{1+z(R)}$. This means that the pressure is in proportion to the curvature of the curve $f(R)$. For an accurate expression of $f(R)$  we propose  a Pade-like form of $f(R)=1/\\sqrt{u}$ with $u\\equiv 1+\\sum\\limits_{n=1}^{N} u_nR^n$. For flat $\\Lambda$ model, the expansion up to $N=7$ has at most an error $< 0.2\\%$  for $z < 1.7$ and all value of $\\Lambda$. We also propose a general method    to determine the equation of state of the universe which has  $< N-1$ free parameters. ",
        "conclusions": "Now if x-matter consists of the scalar field $\\phi$ with the potential $V(\\phi)$, they are related to $\\rho_X$ and $p_X$ as \\ba \\left(\\frac{d\\phi}{dt}\\right)^2&=& \\rho_X+p_X,\\nonumber \\\\ V(\\phi)&=&\\frac{1}{2}(\\rho_X-p_X). \\nonumber \\ea Using $\\rho(R)$ and $p(R)$, we have \\ba &&\\phi-\\phi_0=\\frac{1}{\\sqrt{8\\pi G}}\\int_0^R\\sqrt{-3\\Omega_Mu+\\frac{2}{u}\\frac{d^2u}{dR^2}}dR\\equiv g(R),\\\\ &&V(\\phi)=\\frac{3H_0^2}{\\sqrt{16\\pi G}}\\left(2(\\frac{du}{dR})^2-\\Omega_Mu^3- \\frac{2u}{3}\\frac{d^2u}{dR^2}\\right) \\equiv h(R), \\ea where $\\phi_0$ is the present value of the scalar field . From Eq. (17) we have $R=g^{-1}(\\phi-\\phi_0).$ Then the potential is expressed as $V(\\phi)=h(g^{-1}(\\phi-\\phi_0))$.   In section 3 we assumed that  accurate   $u_n$ for $n=1,\\ldots, N$ are obtained. We here show an example of the determination of $u_n$ for $n=1,\\ldots, 7$. We adopt the redshifts of 38 data with $z > 0.17$ from (\\cite{perl99}). We also adopt the relative error of $R$ for each data as $X_i$. Now let us assume that our universe obeys the $\\Lambda$ model with $\\Omega_M=0.3$. Then we know the theoretical value of $R_i^t$ for each $z_i$. To simulate the real observation, we set $R_i=R_i^t(1\\pm SX_i)$ where S is a scale factor. Let us assume that an accurate observation gives us 38 luminosity distances with $\\sim 0.1\\%$ accuracy  so that the scale factor $S$ is chosen to make the relative error of $R$ for 38 data be $\\sim 0.1\\%$.\\footnote{This might be possible if the statistical error approaches the systematic error in, for example, the SNAP project. } We performed the likelihood analysis for this simulated data and obtained $u_2=0.2240, u_3=0.1505, u_4=0.0585$, $u_5=0.0187, u_6=0.0095$ and $ u_7=0.0053$ with $\\chi^2=36.16$ for 38-7 d.o.f. while theoretical values are  $u_2=0.225, u_3=0.15, u_4=0.054375$, $u_5=0.016875, u_6=0.0068906$ and $ u_7=0.00225$. >From this $\\Omega_M=4/3u_2=0.29866$ or  $\\Omega_M=2u_3=0.3010$ is obtained while $dp/d\\rho=0.0052$ at $u=1$. Assuming the more general equation of state with $w_2=w_3\\ldots=w_ 7=0$, we have $\\Omega_M= 0.31095,  w_0= -1.01783$ and $w_1=-0.04768$. This suggests that we may confirm the $\\Lambda$ model if such accurate luminosity distances are available. We show in Fig. 2 the simulated data and the results of the likelihood analysis. Note that  error bars are extended by a factor 100. The theoretical curve (the dashed line) and the observational curve (the solid line) are almost indistinguishable."
    },
    "0011/astro-ph0011361_arXiv.txt": {
        "abstract": "Re-examination of extensive photometric data of TV~Col reveals evidence for a permanent positive superhump. Its period (6.4~h) is 16 percent 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}  Permanent superhump systems compose a new subclass of cataclysmic variables (CVs), whose existence was established only in the nineties. Systems that belong to this group have quasi-periodicities slightly shifted from their orbital periods, in addition to the binary periods themselves. Unlike SU~UMa systems (see Warner 1995 for a review on SU~UMa systems and CVs in general), which show this behaviour only during superoutbursts, in permanent superhump systems the phenomenon is observed during their normal brightness state. According to Osaki (1996), permanent superhumpers differ from other subclasses of non-magnetic CVs by their 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 `$\\bf positive$ $\\bf superhump$', periodicity which is a few percent larger than the orbital period, is explained as the beat period between the binary motion and the precession of an eccentric disc in the apsidal plane. Periods slightly shorter than the orbital periods have also been seen in several systems. They are known as `$\\bf negative$ $\\bf superhumps$' (Patterson 1999). The observations show a roughly linear relation between the positive superhump period excess as a fraction of the binary period and the  binary period (Stolz \\& Schoembs 1984). Negative superhumps seem to obey a similar rule (Patterson 1999). It has been suggested that negative superhumps are generated by the nodal precession of the accretion disc (Patterson et al. 1993; Patterson 1999); however, there are some theoretical difficulties with this idea (e.g. Murray \\& Armitage 1998).  \\subsection{Periodicities in TV~Col}  The periodicities detected so far in TV~Col and their common interpretations are (Motch 1981; Hutchings et al. 1981; Schrijver, Brinkman \\& van der Woerd 1987; Barrett, O'Donoghue \\& Warner 1988; Hellier, Mason \\& Mittaz 1991; Hellier 1993; Augusteijn et al. 1994):  \\begin{itemize}  \\item 4 day - the nodal precession of the accretion disc  \\item 5.5 hr - the orbital period  \\item 5.2 hr - the negative superhump (the beat between the orbital period and the nodal precession)  \\item 32 min - the spin period  \\end{itemize}  There seems to be a strong connection between positive and negative superhumps. Light curves of many permanent superhumpers show both types of superhumps. In addition, period deficits in negative superhumps are about half period excesses in positive superhumps (Patterson 1999): $\\epsilon_{negative}$$\\approx$--0.5$\\epsilon_{positive}$, where $\\epsilon$=($P_{superhump}$--$P_{orbital})/P_{orbital}$. We, therefore, decided to look in available photometric data on TV~Col for positive superhumps, which would be predicted to have a period near 6.4 h. Here we report the finding of such a periodicity as initially announced by Retter \\& Hellier (2000). ",
        "conclusions": "The photometric data show evidence for a periodicity of 0.265 d in addition to the previously known periods. The repeatability of the peak in three independent datasets makes it 95\\% significant. In addition, the better the data are, the more the period stands out of the noise. Moreover, it has almost exactly the value predicted from the Stolz \\& Schoembs (1984) relation (updated by Patterson 1999) shown in Fig.~2. TV~Col has already been classified as a permanent superhump system because its 5.2-h period was interpreted as a negative superhump. In addition, the new period and the negative superhump obey the relation between the two types of superhumps (Section~1.2). Therefore, the new period is naturally interpreted as a positive superhump.  \\begin{figure}  \\centerline{\\epsfxsize=2.75in\\epsfbox{fig2.eps}}  \\caption{The relation between superhump period excess (over the binary period) and binary period in superhump systems. Empty circles correspond to periods in the SU~UMa systems. Filled circles represent permanent superhumpers. The solid line represents the linear fit to the data. The two tilted dashed lines show the 1-$\\sigma$ error. TV~Col obeys the relation. The upper edge of the period gap (as defined by Diaz \\& Bruch 1997) is marked by the vertical long-dashed line.}    \\end{figure}  According to theory superhumps can appear only in CVs with small mass ratios -- q=$ M_{donor}/M_{compact}$$<$0.33. Hellier (1993) concluded, however, that q=0.62-0.93 from a spectroscopic analysis of the system, but this depended on an interpretation of the emission lines that may not be correct. Using the superhump excess, we find: q=0.95$\\bf \\pm$0.10 -- well above the 0.33 limit suggested by the hydrodynamic simulations, and consistent with the values estimated by Hellier. The mass ratio in TV~Col may thus be above the theoretical limit, perhaps due to its strong magnetic field. Alternatively, TV~Col may be an extreme system, with a very massive white dwarf near the Chandrasekhar mass (1.44$\\bf M_{\\odot}$), and / or an undermassive secondary star."
    },
    "0011/astro-ph0011407_arXiv.txt": {
        "abstract": "We present photometry and high SNR spectroscopy in the classification region of V635~Cas, the optical counterpart to the transient X-ray pulsator 4U\\,0115+63, taken at a time when the circumstellar envelope had disappeared. V635~Cas is classified as a B0.2Ve star at a distance of $7-8$ kpc. We use the physical parameters derived from these observations and the orbit derived from X-ray observations to elaborate a model of the system based on the theory of decretion discs around Be stars. We show that the disc surrounding the Be star must be truncated by the tidal/resonant interaction with the neutron star and cannot be in a steady state. This explains many of the observed properties of 4U\\,0115+63. In particular, because of this effect, under normal circumstances, the neutron star cannot accrete from the disc, which explains the lack of regular Type~I outbursts from the source. ",
        "introduction": "The hard X-ray transient 4U\\,0115+63 (X\\,0115+634) is one of the best studied  Be/X-ray binary systems (see Campana 1996; Negueruela et al. 1997, henceforth N97). More than 50 of these systems, in which a neutron star orbits a Be star in a moderately eccentric orbit, are known (see Negueruela 1998; Bildsten et al. 1997). The Be star is surrounded by a disc of relatively cool material, presumably ejected from the star due to causes unknown, but generally believed to be associated with fast rotation, magnetic fields and/or non-radial pulsations (Slettebak 1988). The presence of the disc gives rise to emission lines in the optical and infrared spectral regions and an excess in the infrared continuum radiation.  The hard X-ray emission is due to the accretion of circumstellar material on to the neutron star companion. Due to their different geometries and the varying physical conditions in the circumstellar disc, Be/X-ray binaries can present very different states of X-ray activity (Stella et al. 1986). In quiescence, they display persistent low-luminosity ($L_{{\\rm x}} \\la 10^{36}$ erg s$^{-1}$) X-ray emission or no detectable emission at all. Occasionally, they show series of periodical (Type I) X-ray outbursts ($L_{{\\rm x}} \\approx 10^{36} - 10^{37}$ erg s$^{-1}$), separated by the orbital period of the neutron star. More rarely, they undergo giant (Type II) X-ray outbursts ($L_{{\\rm x}} \\ga 10^{37}$ erg s$^{-1}$), which do not clearly correlate with the orbital motion. Some systems only display persistent emission, but most of them show outbursts and are termed Be/X-ray transients.    \\begin{figure*} \\begin{picture}(500,250) \\put(0,0){\\special{psfile=fig1.ps angle =0 hoffset=-100 voffset=-240 hscale=120 vscale=120}} \\end{picture} \\caption{The spectrum of V635~Cas in the classification region. Two exposures taken on November 14, 1997, with ISIS on the WHT equipped with the R1200B grating and the EEV10 camera have been combined for this figure. The comparison spectrum is that of the B0.2V standard $\\tau$ Sco. Both spectra have been divided by a spline fit to the continuum for normalisation and smoothed with a $\\sigma = 0.8$\\AA\\ Gaussian function for display.} \\label{fig:bluespec} \\end{figure*}   The transient 4U\\,0115+63 was first reported in the {\\em Uhuru} satellite survey (Giacconi et al. 1972; Forman et al. 1978), though a search of the {\\em Vela 5B} data base revealed that the source had already been observed by this satellite since 1969 (Whitlock et al. 1989). Precise positional determinations by the {\\em SAS 3}, {\\em Ariel V} and {\\em HEAO-1} satellites (Cominsky et al. 1978; Johnston et al. 1978) were used to identify the system with a heavily reddened Be star with a visual magnitude $V \\approx 15.5$ (Johns et al. 1978; Hutchings \\& Crampton 1981), which was subsequently named V635~Cas (Khopolov et al. 1981). Rappaport et al. (1978) used {\\em SAS 3} timing observations to derive the orbital parameters of the binary system, which consists of a fast-rotating ($P_{{\\rm s}} = 3.6 \\: {\\rm s}$) neutron star in a relatively close ($P_{{\\rm orb}} = 24.3 \\: {\\rm d}$) and eccentric ($e = 0.34$) orbit around the Be star (see also Tamura et al. 1992).  Due to the fast rotation of the neutron star, centrifugal inhibition of accretion prevents the onset of X-ray emission unless the ram pressure of accreted material reaches a relatively high value (Stella et al. 1986; N97). The system had only be known to display Type II activity until a short series of Type I outbursts was detected by BATSE and {\\em RXTE} in 1996 (Bildsten et al. 1997; Negueruela et al. 1998). The giant outbursts are associated with large amplitude brightenings of the optical and infrared magnitudes of the counterpart (N97).  This is the first of two papers dedicated to providing a coherent picture of this system and understanding the implications of its unusual behaviour for the general class of Be/X-ray transients. Here we derive the astrophysical parameters of 4U\\,0115+63 and build a model for circumstellar disc around the Be star that allows us to understand the usual quiescent state of the system. In the second paper (Negueruela et al. 2000; henceforth Paper II), we will analyse the temporal evolution of this disc and investigate how its behaviour is connected with the X-ray activity of the source. ",
        "conclusions": "By observing the source during a disc-loss episode, we have been able to determine the stellar parameters of V635~Cas. The object is classified as a B0.2Ve star at a distance of $7-8$ kpc. Both the mass function and the estimated $v\\sin i$ indicate a moderate inclination for the orbital and equatorial planes. The derived distance implies that the source can radiate close to the Eddington luminosity for a neutron star during bright outbursts.  With the newly determined parameters, we have constructed a model for 4U\\,0115+63. Based on the viscous decretion disc model for Be stars, we have numerically solved a criterion for tidal truncation and found that the disc surrounding V635~Cas must be truncated at a resonance radius depending on the viscosity parameter and cannot be in a steady state. Although we have adopted a particular disc model, our conclusion is robust as long as the outflow in the disc is subsonic, a hypothesis supported by several observational facts. Under normal conditions, the neutron star cannot accrete enough material to overcome the centrifugal barrier and switch on the X-ray emission."
    },
    "0011/astro-ph0011541_arXiv.txt": {
        "abstract": "Violation of the axial symmetry of a magnetic field essentially modifies the physics of the plasma outflow in the magnetosphere of rotating objects. In comparison to the axisymmetric outflow, two new effects appear: more efficient magnetocentrifugal acceleration of the plasma along particular field lines and generation of  MHD waves. Here, we use an ideal MHD approximation to study the dynamics of a cold wind in the nonaxisymmetric magnetosphere. We obtain a self-consistent analytical solution of the problem of cold plasma outflow from a slowly rotating star with a slightly nonaxisymmetric magnetic field using perturbation theory. In the axisymmetric (monopole-like) magnetic field, the first term in the expansion of the terminating energy of the plasma in powers of $\\Omega$ is proportional to $\\Omega^4$, where $\\Omega$ is the angular velocity of the central source. Violation of the axial symmetry of the magnetic field crucially changes this dependence. The first correction to the energy of the plasma becomes proportional to $\\Omega$. Efficient magnetocentrifugal acceleration occurs along the field lines curved initially in the direction of the rotation. I argue  that  all the necessary conditions for the efficient magnetocentrifugal acceleration of the plasma exist in the radio pulsar magnetosphere. We calculated  the first correction of the rotational losses due to the generation of the MHD waves  and analysed  the plasma acceleration by these waves. ",
        "introduction": "The outflow of plasma from rotating magnetised objects is a widespread phenomenon in the Universe. It is observed in a wide range of astrophysical objects of different natures and scales. Some of these objects eject relativistic plasma. The Lorentz factor of plasma ejected by microquasars and AGN is of the order 10 (Mirabel \\& Rodrigues \\cite{mirabel}, Begelman et al. \\cite{bbr}, Pelletier et al. \\cite{pelletier}), although there is evidence that in Blazars the plasma is accelerated to the Lorentz factor $\\sim 10^6$ (Aharonian et al. \\cite{ah}). Radio pulsars also accelerate the plasma to the Lorentz factor $\\sim 10^6$ (Kennel \\& Coroniti \\cite{kennel}, Coroniti \\cite{coroniti}, Arons \\cite{arons96}). The problem of the acceleration of the plasma to such high energies  remains one of the most important  unsolved problems in modern relativistic astrophysics.  The acceleration of plasma can occur in principle due to different dissipative as well as nondissipative processes. Hypothesis have been proposed  about the singular behaviour of the relativistic plasma outflow at the ``light surface'' (\\cite{beskin83}) or at the light cylinder (Mestel \\& Shibata \\cite{ms}) of radio pulsars. These dissipative structures could provide  the  plasma acceleration. However, recent direct  numerical calculations of the relativistic outflows in MHD approximation (Bogovalov \\cite{bog97}) and force-free approximation (Contpoulos et al. \\cite{janis}) have demonstrated that the relativistic plasma flow is actually regular everywhere.   Another mechanism of  acceleration due to dissipative processes was proposed by Michel (1982) and Coroniti (1990) in application to radio pulsars. Michel was the first who noticed that the number of charge carriers in the wind from radio pulsars  is not enough to sustain the necessary return current in the wind with a striped magnetic field formed at the outflow from the oblique rotator.  Coroniti (1990) considered the reconnection process in this magnetic field and came to the conclusion that it can provide the necessary acceleration of the winds from radio pulsars. Up to now this process was considered as the only possible mechanism explaining pulsar wind acceleration. However, recently Luybarskii \\& Kirk (2000) have revisited  the process of wind acceleration due to  magnetic field reconnection. They  concluded that well before the winds from pulsars such as the Crab pulsar reach high Lorentz factors they will be terminated by the interstellar medium.   A lot of efforts has been spent to obtain the plasma acceleration in frameworks of ideal MHD. There is a general opinion that in all astrophysical objects ejecting plasma, the magnetic field plays an essential role in the dynamics of the plasma. In particular, the magnetic field of a rotating object itself can provide the acceleration of the plasma due to a magnetocentrifugal (below MC) mechanism (Blandford \\& Payne \\cite{bp}). The ideal MHD outflows of plasma in different models  and different approximations were investigated in a series of works (Michel \\cite{michel69}, Sulkanen \\& Lovelace \\cite{lovlace}, Ferreira \\cite{ferreira}, Vlahakis \\& Tsinganos \\cite{tsin},  Bogovalov \\& Tsinganos \\cite{bogtsin}, Krasnopolsky  et al. \\cite{kb}, Ustyugova et al. \\cite{ustyug}). It follows from these studies that nonrelativistic plasma can be accelerated rather efficiently due to this mechanism. As far as relativistic plasma is concerned,  the efficiency of the acceleration of this plasma  appears insufficient for astrophysical applications. There were hopes that divergence of the magnetic flux tubes somewhere beyond the  light cylinder will result in more efficient acceleration of the plasma due to the pressure gradient of the toroidal magnetic field (Begelman \\& Li \\cite{bl}, Takahashi \\& Shibata \\cite{takshib}). The only known mechanism for this magnetic flux tube divergence  is  magnetic self-collimation of the magnetised winds (Heyvaerts \\& Norman \\cite{heyverts}, Chiueh et al. \\cite{chuk}). This mechanism can, in principle, provide divergence of the field lines at the equatorial plane such that the poloidal field decreases with distance faster than $r^{-2}$. This divergence can take place in the limiting range of distances from the source, since finally the wind  expands radially near the equator at $r\\rightarrow \\infty$  (Heyvaerts \\& Norman \\cite{heyverts}, Chiueh et al. \\cite{chuk}).  Numerical self-consistent solution of the problem of steady-state  plasma outflow in a model with the an initially monopole-like magnetic field shows that collimation really increases the efficiency of nonrelativistic plasma acceleration at the equator (Bogovalov \\& Tsinganos \\cite{bogtsin}). However, collimation of the  relativistic plasma outflow  appears negligibly small at large Lorentz factors  (Bogovalov \\cite{bog97}, Bogovalov \\cite{bog00}). Therefore, the relativistic plasma is practically not accelerated by the pressure gradient of the toroidal magnetic field. This conclusion is valid not only for the monopole-like model.  The acceleration of the wind due to the pressure gradient of the toroidal magnetic field occurs at the large distances compared to the dimension of the central source (Begelman \\& Li \\cite{bl}, Takahashi \\& Shibata \\cite{takshib}). The stationary wind from any source is monopole-like at these distances. Therefore, the poloidal magnetic field  also becomes monopole-like, since the magnetic field is frozen into the plasma. Thus, the solutions obtained in the monopole-like model actually describe general properties of any axisymmetric wind at large distances from the source.  The application of the  monopole-like model is not limited by large distances. This model  also gives  upper  limits  on  the  efficiency of  the MC acceleration of the plasma in any other axisymmetric  magnetic field (dipolar for example). It is qualitatively clear that the slower the magnetic field decreases with  distance, the more efficiently it accelerates the plasma due to a MC mechanism. The monopole-like magnetic field  falls down  as  $r^{-2}$, while the dipolar magnetic field falls down faster than $r^{-2}$. Correspondingly, the plasma is expected to be more efficiently accelerated in the monopole-like magnetic field in compare with the acceleration in the dipolar one. This is why the results obtained in the monopole-like model rule out efficient MC acceleration of the plasma in any other more realistic axisymmetric model. Calculations  performed by Contopoulos et al. (\\cite{janis}) confirm this conclusion.  Thus, up to now all the attempts to propose an effective mechanism of relativistic plasma acceleration have not been successful. That is why the search for new possible mechanisms to explain relativistic plasma acceleration remains one of the most important problems of  relativistic astrophysics.   The disappointing conclusion regarding the low efficiency of the MC acceleration  is valid only for   axisymmetric magnetic fields, as previous studies considered  this process only within the frameworks of  axisymmetric models. However,  the magnetic field of  real astrophysical objects is not axisymmetric. Radio pulsars give us the  brightest example of this relation. Therefore, it is important to know how  the non axisymmetry of the magnetic field affects the process of  MC acceleration. In this paper we try to answer this question.  To clarify the role of the nonaxisymmetry of the magnetic field for the MC  acceleration  of    plasma,  a  model  with  the initially monopole-like magnetic field is used in this paper. This model is often used to investigate the processes of plasma collimation and acceleration  in the rotating magnetosphere (Michel \\cite{michel69}, Mestel \\& Selley \\cite{mestel}, Sakurai  \\cite{sakurai},   Bogovalov   \\cite{bog92}, Beskin et al. \\cite{vasia98}, Bogovalov  \\&  Tsinganos  \\cite{bogtsin}, Tsinganos   \\& Bogovalov \\cite{tsbog}). The word ``initially\" implies that the magnetic field of the non rotating star looks like the magnetic field of the ``magnetic monopole\". This model is very convenient since there are no closed field lines in the flow. It remarkably simplifies the analysis of  the plasma acceleration.  The main goal of this paper is to  answer  the question: how does the nonaxisymmetry of the magnetic field of a star affect the process of  plasma acceleration? An attempt to solve the problem of the structure of the magnetosphere of an oblique rotator with a dipole magnetic field in mass-less approximation has been done by Beskin et al. (\\cite{beskin83}). The first step in the solution of this problem in the MHD approximation has been done in the work of Bogovalov (\\cite{bog99}). In this work, the problem of plasma flow in the magnetosphere of the oblique rotator with an initially split-monopole magnetic field was solved. The modulus of the magnetic field was axisymmetric and only the direction of the magnetic field varied with time and the azimuthal angle. It was found that the acceleration remains non efficient in this flow as well. In the present paper, we are interested in the affect of the nonaxisymmetry of the  modulus of the magnetic field on the process of  plasma acceleration.  Plasma flow is described by a system of non linear equations in partial derivatives. Usually  these equations can be integrated only numerically. There is one exception from this rule.  Sometimes the problem can be solved self-consistently and analytically if we are interested in small corrections to the known solution. These corrections can arise if the known flow is perturbed slightly by additional  small forces. In particular,  slow rotation can be considered as a small perturbation of some known flow from the non rotating star. This approach was firstly successfully applied for the numerical self-consistent solution of the flow of the solar wind by Nurney and Suess (\\cite{ns}). Later Bogovalov (\\cite{bog92}) has used this approach for the fully analytical, self-consistent solution of the problem of the cold plasma outflow, relativistic as well as nonrelativistic, from a slowly rotating star. In recent years, this method has also been used to solve  a range of problems by Beskin (see Beskin \\& Okamoto \\cite{bsknokmto}) and references therein). The idea of the present work is the following: if the nonaxisymmmetry of the magnetic field modifies the process of MC acceleration of the plasma, this modification can be seen at the level of the slow rotation of the object. Therefore, to answer the question: does  the nonaxisymmetry of the magnetic field modify the process of  acceleration of the plasma, we investigate the acceleration of  plasma in the nonaxisymmetric magnetic field at slow rotation.   This paper is organised as follows. In Sect. 2  we present ideal MHD equations defining the dynamics of cold relativistic plasma. A self-consistent analytical solution of the problem is given in Sec. 3. The solution in the wave zone is given in Sect. 4. The possible astrophysical implications of the results are  discussed in Sec. 5. ",
        "conclusions": ""
    },
    "0011/astro-ph0011294_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the 19th century the American mathematician G.W. Hill devised a simple and useful approximation for the motion of the moon around the earth with perturbations by the sun.  To most dynamical astronomers ``Hill's Problem''\\index{Hill's problem} still means a model for motions in the solar system in which two nearby bodies move in nearly circular orbits about another much larger body at a great distance.  These lectures have, however, been motivated by a problem in stellar dynamics\\index{stellar dynamics}.  Consider a star in a star cluster\\index{star cluster} which is itself in orbit about a galaxy (Figure \\ref{derivation}).  The star, cluster and galaxy take the place of the moon, earth and sun, respectively.  The potentials of the cluster and galaxy are not those of a point mass, and the galactic orbits of the star and cluster may be far from circular.  Nevertheless Hill's problem is a good starting point, and it can be modified easily to accommodate the differences.  In section 2 we outline a derivation of Hill's equations, and in section 3 we summarise the appropriate extensions.  Stars gradually escape\\index{escape} from star clusters. This has been expected on theoretical grounds for many years, ever since a paper by Ambartsumian (1938). Recently, deep observations have confirmed this (e.g. Leon et al 2000), by revealing faint streams of stars around a number of the globular clusters of our Galaxy.  Loosely speaking we can say that a star can only escape if its energy exceeds some critical energy.  The energies of stars change slightly as a result of two-body gravitational encounters within clusters, though the time scale on which this happens (the {\\sl relaxation time\\index{relaxation time} scale}) is very long, of order $10^9$yr.  But the orbital motions of stars within clusters have much smaller time scales of order $10^6$yr, and until recently it was thought that escaping stars would leave on a similar time scale.  With this assumption, relaxation is the bottleneck, and so the escape time scale (e.g. the time taken for half the stars to escape) should vary with the relaxation time.  Nowadays it is possible to simulate the evolution of modest-sized star clusters with $3\\times10^4$ or more members, and the predicted escape time scale can be checked empirically.  Unfortunately the results contradict the theory (Figure {\\ref{collaborative}}).  As these simulations\\index{$N$-body simulations} require considerable extrapolation in particle number $N$ to be applicable to real clusters (for which $N\\sim 10^6$) the error of the theory is serious.  It turns out that the assumption of rapid escape is the main source of error (Fukushige \\& Heggie 2000, Baumgardt 2000a,b).  In fact some stars above the escape energy never  escape\\index{escape} (unless some other dynamical process comes into play), and others take much longer to escape than had been generally thought.   \\begin{figure} [htbp] \\centering \\includegraphics[width=10cm,clip,trim=0 0 0 0,angle=-90]{th.eps} \\caption{Results of numerical experiments (Aarseth \\& Heggie, unpublished) on the escape of stars from star clusters\\index{star cluster}.  The time for half the stars to escape is plotted against the original number $N$ of stars in the simulation.  Points are averages over several simulations\\index{$N$-body simulations} at each $N$, except the largest value.  The continuous line shows the prediction of theory, i.e. proportional to the relaxation time\\index{relaxation time} (see text), and the dashed line is an empirical fit.} \\label{collaborative} \\end{figure}  With this motivation, the remaining sections of these lectures are devoted to the dynamics of escape.  Section 4 analyses the very definition of escape, which is not as straightforward as in more familiar situations.   The last two sections show some ways in which the computation of the escape rate can be approached.  The main result of section 5 concerns the way in which the time scale of escape depends on the energy, and outlines how this resolves the problem of Figure {\\ref{collaborative}}.   Much more difficult, from a theoretical point of view, is determining the {\\sl distribution} of escape times, and some relevant ideas are introduced in Section 6. ",
        "conclusions": ""
    },
    "0011/astro-ph0011206_arXiv.txt": {
        "abstract": "After decades of one-dimensional nucleosynthesis calculations, the growth of computational resources has meanwhile reached a level, which for the first time allows astrophysicists to consider performing routinely realistic multidimensional nucleosynthesis calculations in explosive and, to some extent, also in non-explosive environments.  In the present contribution we attempt to give a short overview of the physical and numerical problems which are encountered in these simulations. In addition, we assess the accuracy that can be currently achieved in the computation of nucleosynthetic yields, using multidimensional simulations of core collapse supernovae as an example. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011030_arXiv.txt": {
        "abstract": "Associated absorption lines (AALs) are valuable probes of the gaseous environments near quasars. Here we discuss high-resolution (6.7 \\kms ) spectra of the AALs in the radio-loud quasar 3C 191 (redshift $z=1.956$). The measured AALs have ionizations ranging from \\ion{Mg}{1} to \\ion{N}{5}, and multi-component profiles that are blueshifted by $\\sim$400 to $\\sim$1400~\\kms\\ relative to the quasar's broad emission lines. These data yield the following new results.  1) The strengths of excited-state \\ion{Si}{2}$^*$ AALs indicate a density of $\\sim$300 \\pcc\\ in the Si$^+$ gas.  2) If the gas is photoionized, this density implies a distance of $\\sim$28 kpc from the quasar. Several arguments suggest that all of the lines form at approximately this distance.  3) The characteristic flow time from the quasar is thus $\\sim$$3\\times 10^7$ yr.  4) Strong \\ion{Mg}{1} AALs identify neutral gas with very low ionization parameter and high density. We estimate $n_H\\ga 5\\times 10^4$ \\pcc\\ in this region, compared to $\\sim$15 \\pcc\\ where the \\ion{N}{5} lines form.  5) The total column density is $N_{\\rm H}\\la 4\\times 10^{18}$ \\cmsq\\ in the neutral gas and $N_{\\rm H}\\sim 2\\times 10^{20}$ \\cmsq\\ in the moderately ionized regions. These column densities are consistent with 3C 191's strong soft X-ray flux and the implied absence of soft X-ray absorption.  6) The total mass in the AAL outflow is $M\\sim 2\\times 10^9$ \\Msun , assuming a global covering factor (as viewed from the quasar) of $\\sim$10\\% .  7) The absorbing gas only partially covers the background light source(s) along our line(s) of sight, requiring absorption in small clouds or filaments $<$0.01 pc across. The ratio $N_H/n_H$ implies that the clouds have radial (line-of-sight) thicknesses $\\la$0.2 pc.  These properties might characterize a sub-class of AALs that are physically related to quasars but form at large distances. We propose a model for the absorber in which pockets of dense neutral gas are surrounded by larger clouds of generally lower density and higher ionization. This outflowing material might be leftover from a blowout associated with a nuclear starburst, the onset of quasar activity or a past broad absorption line (BAL) wind phase. ",
        "introduction": "Associated absorption lines (AALs) are important diagnostics of the gaseous environments of quasars and active galactic nuclei (AGNs). The lines are defined empirically as having velocity widths less than a few hundred \\kms\\ and absorption redshifts, $z_a$, within a few thousand \\kms\\ of the quasar's emission-line redshift, $z_e$ (Weymann \\etal 1979, Foltz \\etal 1986). The requirement for \\zaz\\ makes AALs more likely to be physically related to the quasars than are the other narrow absorption lines (at \\zllz ) in quasar spectra  (see Rauch 1998 for a review of the unrelated \\zllz\\ systems). The narrow velocity widths further distinguish AALs from the class of broad absorption lines (BALs), whose widths and maximum blueshifted velocities both typically exceed 10,000 \\kms\\ (Weymann \\etal 1991). BALs clearly form in high-velocity winds from quasar engines (e.g. Turnshek 1988), but AALs can form potentially in a variety of environments --- ranging from energetic outflows like the BALs to relatively quiescent gas at large galactic or inter-galactic distances (see Hamann \\& Brandt 2000 for a general review, also Tripp \\etal 1996, Hamann \\etal 1997a, Barlow \\& Sargent 1997, Barlow, Hamann \\& Sargent 1997). More work is needed to locate individual AAL absorbers, quantify their kinematic and physical properties, and understand the role of the AGNs and/or host galaxies in providing the source of their material and kinetic energy.  One interesting property is that, among radio-loud quasars, AALs appear more frequently and with greater strength in sources with ``steep'' radio spectra and/or lobe-dominated radio morphologies (Wills \\etal 1995, Barthel \\etal 1997, Richards \\etal 2000, also Brotherton \\etal 1998 and references therein). This weak correlation is usually attributed to an orientation effect, whereby AAL regions reside preferentially near a disk or torus that is aligned perpendicular to the radio jet axis (but see Richards \\etal 2000 for an alternative interpretation). However, it is not clear if this (possible) AAL geometry has its origins on small scales related to the inner black hole/accretion disk, or on much larger scales related to the host galaxy. It is also not clear if a disk-like geometry applies as well to the AALs in radio-quiet sources.  3C 191 (Q0802+103, $z_e = 1.956$) is a radio-loud quasar having both strong AALs (Burbidge, Lynds \\& Burbidge 1966, Stockton \\& Lynds 1966, Williams \\etal 1975, Anderson \\etal 1987) and a bipolar, lobe-dominated radio structure (Akujor et al. 1994). It therefore follows the AAL--radio morphology correlation noted above. 3C 191 also provides a rare opportunity to define the distance between the quasar and the absorbing gas because several of its AALs arise from excited energy states, e.g. \\ion{C}{2}$^*$ \\lam 1336 and \\ion{Si}{2}$^*$ \\lam 1265,1533 (Bahcall, Sargent \\& Schmidt 1967, Williams \\etal 1975). The strengths of the excited-state lines, compared to their resonant counterparts, \\ion{C}{2} \\lam 1335, \\ion{Si}{2} \\lam 1260,1527, provide measures of the gas density needed to populate the upper levels. The density in turn constrains the absorber's distance from the quasar, with the reasonable assumption that the gas is in photoioization equilibrium with the quasar radiation field.  Williams et al. (1975) already estimated a density of $n_e\\sim 1000$ \\pcc\\ and a radial distance of $R\\sim 10$ kpc for the AAL region in 3C 191. We reobserved 3C 191 with higher spectral resolution and wider wavelength coverage to 1) obtain more reliable densities as a function of velocity, 2) search for \\ion{Fe}{2}$^*$ AALs, which might be revealing of a much higher density environment (Wampler \\etal 1995, Halpern \\etal 1996), 3) revisit the question of this absorber's origin and location, and 4) obtain better constraints on the AAL region dynamics, abundances and overall physical structure.  Sections 2 and 3 below describe the observations and results. Section 4 provides measurements and analyses of the AALs. Section 5 draws further inferences and discusses physical models. Throughout this paper we define solar abundances by the meteoritic results in Grevesse \\& Anders (1989), and we use atomic transition data from the compilation by Verner, Verner \\& Ferland (1996). ",
        "conclusions": "\\subsection{Implications for the Structure and Dynamics}  The results in \\S3 and \\S4 lead to following important conclusions regarding the AAL environment.  1) We have already noted that the absorber has a complex velocity structure that appears qualitatively similar in all lines and ions (\\S4.1). Stronger AALs, e.g. of higher ionization, are more smoothly distributed in velocity, but it seems clear that all of the lines trace the same overall physical structure. We conclude that the radial distance $R\\approx 28$ kpc derived for the Si$^+$ zone (\\S4.6) should be roughly characteristic of the entire AAL region (see also item 4 below).  2) The AAL gas appears to be outflowing from the quasar at velocities from $\\sim$400 to $\\sim$1400 \\kms\\ (\\S4.1). If the gas is not accelerating (or decelerating), then the time scale for this outflow reaching radius $R$ is \\begin{equation} t\\ \\approx\\ 3\\times 10^7\\ \\left({{R}\\over{28\\, {\\rm kpc}}}\\right) \\left({{1000\\, {\\rm km\\, s}^{-1}}\\over{v}}\\right)~~{\\rm yr} \\end{equation} where $v$ is a characteristic velocity. The overall AAL region will acquire a radial thickness during this expansion because different gas components move at different speeds. In 3C 191, the radial thickness, $\\Delta R$, could be comparable to the radius, $R\\approx 28$ kpc, because the velocity dispersion implied by the line widths, $\\Delta v \\approx 1000$ \\kms , is similar to the average flow speed.  3) The significant presence of \\ion{Mg}{1}, compared to \\ion{Mg}{2}, identifies a neutral gas component whose survival cannot be attributed to radiative shielding downstream from an \\ion{H}{2}--\\ion{H}{1} recombination front (\\S4.4). The \\ion{Mg}{1} region must have a low ionization parameter and therefore a high density. If we adopt $n_H\\approx 300$ \\pcc\\ and $\\log U\\approx -2.8$ for the \\ion{Si}{2} region (\\S4.6), then simple scaling based on $\\log U\\la -5$ in the neutral gas (\\S4.4) implies $n_H\\ga 5\\times 10^4$ \\pcc\\ in that region. In contrast, $\\log U\\sim -1.5$ in the \\ion{N}{5} zone indicates $n_H\\sim 15$ \\pcc\\ there.  4) The \\ion{Si}{2} \\lam 1527 and \\ion{Si}{2}$^*$ \\lam 1533 line profiles are surprisingly similar (Fig. 2) given the wide range of densities present. The nearly constant line ratio indicates a nearly constant density (within a factor of $\\sim$2) at all velocities. This result could be caused by a selection bias: the absorber may have a wide range of densities at each velocity, but the specific ionization requirements of Si$^+$ might persistently lead to just a narrow range in $U$ and therefore $n_{\\rm H}$ controlling the \\ion{Si}{2} and \\ion{Si}{2}$^*$ line strengths. In any case, the similar $n_{\\rm H}$ values inferred for different \\ion{Si}{2} velocity components support the argument (item 1 above) that the AAL region does not span a wide range in radial distance from the quasar.  5) Partial line-of-sight coverage of the quasar emission sources (\\S4.2) implies that the AAL clouds have characteristic sizes similar to or smaller than the projected area of the emitters. The \\Lya\\ and \\ion{C}{4} absorption lines sit atop strong BELs (Fig. 2); they might fully cover the continuum source while partly covering the larger BEL region. The characteristic size of these absorbers might therefore be as large as 0.1 -- 1 pc (i.e. the size of the BEL region, e.g. Peterson 1993, Kaspi \\etal 2000). However, the partial coverage inferred from the \\ion{Mg}{2}, \\ion{Al}{3} and \\ion{Fe}{2} lines must involve the continuum source, which requires absorber size scales $<$0.01 pc (for standard accretion disk models of the continuum emission, Netzer 1992).  6) Another constraint on the size scales comes from the ratio $N_H/n_H$. If the absorbing gas completely fills the volume it encompasses, then the radial thickness of the entire absorber would be of order, \\begin{equation} \\Delta R\\ \\approx\\ 0.2\\, \\left({{N_H}\\over{2\\times 10^{20}\\, {\\rm cm}^{-2}}}\\right) \\left({{300\\, {\\rm cm}^{-3}}\\over{n_H}}\\right) ~~{\\rm pc} \\end{equation} using parameters from \\S4.5 and \\S4.6. However, if the absorber is composed of discrete clouds that fill only part of the encompassed volume, then the overall AAL region could have a much greater radial thickness (item 2 above) while the individuals clouds are small compared to $\\Delta R$ in Equation 11.  7) The mass of the AAL region, $M$, depends on its radial distance, total column density and {\\it global} covering factor, $Q\\equiv \\Omega /4\\pi$ (where $\\Omega$ is the solid angle subtended by the absorber as ``seen'' from the central quasar). The value of $Q$ is not known, but it is not likely to exceed the detection frequency of AALs among radio-loud quasars\\footnote{The detection frequency sets an approximate upper limit on $Q$ because some fraction of the systems counted as AALs will have a different physical origin than the absorber in 3C 191 (\\S1).}, e.g. $\\sim$30\\% (G. Richards, private communication). If we let $Q_{0.1}$ represent the covering factor relative to $Q=0.1$, the total mass is given by $$M\\ \\approx\\ 2\\times 10^{9} ~ Q_{0.1} \\left({{\\mu_H}\\over{1.4}}\\right) \\left({{N_H}\\over{2\\times 10^{20}\\, {\\rm cm}^{-2}}}\\right) $$ \\begin{equation} \\times\\ \\left({{R}\\over{28\\, {\\rm kpc}}}\\right)^2 ~{\\rm M}_{\\odot} \\end{equation} where $\\mu_H$ is the mean molecular weight per H particle. The total kinetic energy in the AAL outflow is therefore, \\begin{equation} K\\ \\approx\\ 2\\times 10^{58}\\,\\left({{M}\\over{3\\times 10^9\\, \\Msun}}\\right) \\left({{v}\\over{1000\\, {\\rm km\\, s}^{-1}}}\\right)^2 ~ {\\rm ergs} \\end{equation} For comparison, the much faster ($v\\sim 10^4$ \\kms ) but smaller-scale ($R\\sim 0.1$ pc) BAL outflows observed in other quasars are believed to contain total masses (at any instant) of $\\sim$1--10 \\Msun\\ for $Q_{0.1}\\sim 1$ to 3. Over a quasar's lifetime, say 10$^8$ yr (Haehnelt, Natarajan \\& Rees 1998), a BAL wind could eject a total of $\\sim$10$^7$--10$^8$ \\Msun\\ with $K \\approx\\ 10^{58}$--10$^{59}$ ergs (Hamann \\& Brandt 2000).  8) Emission lines from the AAL gas could be quite strong, depending on the actual values of $U$ or $N_{\\rm H}$. For example, we calculated emission line strengths for a photoionized plasma consistent with the \\ion{Si}{2} AAL region described above, namely, having solar abundances, an incident spectrum as defined in \\S4.4, and the following physical paramters: $R\\approx 28$ kpc, $\\log U\\approx -2.8$, $N_{\\rm H} = 2\\times 10^{20}$ \\cmsq , and $n_H\\approx 300$ \\pcc . The strongest predicted emission lines within our wavelength coverage for 3C 191 are \\Lya\\ and \\ion{Mg}{2} \\lam 2799, with  rest-frame equivalent widths of $\\sim$44$Q_{0.1}$ \\AA\\ and $\\sim$4$Q_{0.1}$ \\AA . These results should be compared to the rest-frame equivalent widths of the measured BELs, e.g. $\\sim$120 \\AA\\ in \\Lya\\ and $\\sim$17 \\AA\\ in \\ion{Mg}{2} (as estimated without the superposed AALs). The predicted emission lines could therefore be present but ``hidden'' in the measured BELs. Clearly, however, $Q_{0.1}\\sim 1$ is close to an upper limit by this test. Another constraint is that $Q_{0.1}>1$ could lead to too much line emission ``filling in'' the bottoms of the AAL troughs (depending on the global line-of-sight velocity distribution of the emitting gas; see Hamann, Korista \\& Morris 1993 and Hamann \\& Korista 1996 for similar arguments related to BALs). Future observations at longer wavelengths might provide more stringent upper limits on $Q_{0.1}$. In particular, our calculations predict that the strongest lines in the rest-frame near-UV/visible should be H$\\alpha$, H$\\beta$, the \\ion{O}{2} \\lam 3727 doublet, and \\ion{O}{3} \\lam 5007, with rest equivalent widths of $\\sim$$55Q_{0.1}$ \\AA , $\\sim$11$Q_{0.1}$ \\AA , $\\sim$22$Q_{0.1}$ \\AA , and $\\sim$88$Q_{0.1}$ \\AA , respectively.  \\subsection{Toward a Physical Model}  Figure 6 shows a highly schematic model of the AAL region, wherein pockets of dense neutral gas are surrounded by a diffuse, spatially distributed medium of generally higher ionization. The diffuse clouds contain most of the total column density (\\S4.5). Their greater size and/or greater numbers lead to more complete coverage in both velocity and projected area (for example, in the \\ion{C}{4} doublet), compared to the lower-ionization gas (e.g. \\ion{Mg}{1} \\lam 2853).  Note that the extended regions cannot have exclusively high ionization levels because the measured line-of-sight coverage fractions do not correlate simply with ionization. For example, low ionization lines can have either high (e.g. \\ion{C}{2} \\lam 1334, \\ion{Si}{2} \\lam 1260) or low (\\ion{Fe}{2}) coverage fractions (\\S4.2). The amount of coverage in both space and velocity must depend at least partly on the line's oscillator strength. In other words, the coverage fraction scales with the line optical depth. Stronger (more optically thick) lines have greater contributions from the diffuse extended gas, resulting in greater coverage, whereas weak lines sample mainly the higher column density material in more compact regions.  A major concern with any cloud model is the cloud survival. The lifetime of a cloud without pressure confinement is of order the sound-crossing time. For a nominal temperature of $10^4$ K and a maximum cloud size of $\\Delta R \\la\\ 0.2$ pc (Eqn. 11), the cloud survival time, $\\la$$2\\times 10^4$ yr, is much less than the characteristic flow time, $t\\approx\\ 3\\times 10^7$ yr (Eqn 10). Therefore pressure confinement appears necessary. The problem of understanding this confinement has plagued cloud models of both the BAL and BEL regions of AGNs. Possible solutions include external pressure from a magnetic field or a surrounding hot, low-density (and transparent) plasma (Weymann, Turnshek \\& Christiansen 1985, Arav, Li \\& Begelman 1994, Emmering, Blandford \\& Shlosman 1992, DeKool 1997, Feldmeier \\etal 1997).  The key remaining question is, what provides the source of material and kinetic energy for the AAL outflow? The flow time ($\\sim$$3\\times 10^7$ yr) is comparable to predicted quasar lifetimes (Terlevich \\& Boyle 1993, Haehnelt et al. 1998). It therefore seems likely that the AAL gas originated much nearer the quasar, perhaps coincident with the onset of quasar activity.  We have already noted (\\S5.1, item 7) that the kinetic energy in this AAL region is comparable to the typical energy in BAL winds (integrated over a quasar's lifetime). Therefore, quasars are capable of driving winds with this total energy. In addition, most current models of BAL winds have them preferentially located near the plane of the accretion disk (Murray \\& Chiang 1995, Emmering \\etal 1992, Wills, Brandt \\& Laor 1999 and references therein). This geometry is reminiscent of the equatorial structure (tentatively) inferred for AAL regions in radio-loud quasars (\\S1).  However, it is unlikely that the AAL outflow in 3C 191 is simply an extended remnant of a BAL wind because 1) the terminal velocity of a BAL wind should be of order 10$^4$ \\kms\\ instead of 1000 \\kms , and 2) the total mass in AAL gas is at least an order of magnitude larger than expected for BAL winds (item 7 in \\S5.1). To produce the observed AALs, a high-velocity BAL-like wind would have to be decelerated by interaction with ambient galactic material and then, probably, entrain some of that material (to add mass) along the way.  An alternative possibility is that the AALs form in gas that was expelled by stellar processes, e.g. in a galactic ``superwind'' as observed in low-redshift starburst galaxies (Heckman, Armus \\& Miley 1990, Heckman \\etal 2000). The sizes, masses, velocities, etc. inferred for superwinds in luminous starbursts are consistent with our estimates of these quantities in 3C 191. The superwinds also contain cool dense clouds (giving rise to \\ion{Na}{1} absorption lines) embedded in a hot ($\\sim$10$^7$ K) X-ray emitting plasma (see also Heckman \\etal 1996 and references therein). If a superwind model does apply to 3C 191, the characteristic flow time of the AAL gas ($\\sim$$3\\times 10^7$ yr) might represent the time elapsed since the starburst episode.  No matter what scenario accounts for the AALs in 3C 191, it is important to keep in mind that AALs in different objects can probe very different physical phenomena. For example, AALs in other quasars often have higher blueshifted velocities than those in 3C 191. Some narrow absorption lines have blueshifts above the arbitrary 5000 \\kms\\ AAL threshold, even though there is strong evidence for their being  intrinsic to the quasar environments (Hamann \\etal 1997a and 1997b, Barlow \\& Sargent 1997, Barlow, Hamann \\& Sargent 1997, Richards 2000). A galactic superwind certainly cannot explain these  high-velocity absorbers. Most AAL systems also do not have low-ionization lines like 3C 191 (e.g. Junkkarinen, Hewitt \\& Burbidge 1991, Hamann 1997). It is possible that low ionization AALs, which allow us to locate the absorber via \\ion{Si}{2}$^*$, select in favor of large absorber--quasar distances. In particular, all of the known AAL systems with these excited-state lines have distances $\\ga$10 kpc (e.g. Barlow \\etal 1997, Tripp, Lu \\& Savage 1996, Morris \\etal 1986, Sargent, Boksenberg \\& Young 1982). Other AALs are known to form much closer to the quasars, possibly within a few pc in outflows similar to the BALs (Hamann \\etal 1997a and 1997b, Barlow \\& Sargent 1997, Barlow \\etal 1997).  Given this diversity, it is interesting to note that 3C~191 does not follow the trend identified by Brandt \\etal (2000) for small X-ray to UV continuum flux ratios accompanying strong \\ion{C}{4} absorption equivalent widths (\\S4.5). That correlation nominally points to a relationship between the strength of the AALs and the strength of continuous (bound-free) absorption in X-rays. BAL quasars are at one extreme in this relationship --- having both strong UV lines and strong absorption in X-rays (see also Green \\& Mathur 1996, Gallagher \\etal 1999). 3C 191 might contain a different class of absorber (e.g. much farther from the active nucleus) than the majority of sources discussed by Brandt \\etal (2000)."
    },
    "0011/astro-ph0011389_arXiv.txt": {
        "abstract": "We present a comprehensive study of the spectrum of the narrow-line Seyfert 1 galaxy \\rej, summarizing the information obtained from the optical to X-rays with observations from the William Herschel 4.2m Telescope (WHT), the Hubble Space Telescope (\\hst), the Extreme UltraViolet Explorer (\\euve), \\rosat, \\asca\\ and \\sax. The \\sax\\ spectra reveal a soft component which is well-represented by two blackbodies with $kT_{\\rm eff}$=60~\\eV\\ and 160~\\eV, mimicking that expected from a hot, optically-thick accretion disc around a low-mass black hole. This is borne out by our modeling of the optical to X-ray nuclear continuum, which constrains the physical parameters of a NLS1 for the first time. The models demonstrate that \\rej\\ is likely to be a system with a nearly edge-on accretion disk (60 to 75$^\\circ$ from the disk axis), accreting at nearly Eddington rates (0.3 to 0.7 $L_{\\rm Edd}$) onto a low mass black hole (M$_{bh}\\sim$2 to 10$\\times 10^6$ M$_{\\odot}$). This is consistent with the hypothesis that NLS1s are Seyfert-scale analogies of Galactic Black Hole Candidates. The unusually high temperature of the big blue bump reveals a flat power-law like continuum in the optical/UV which is consistent with an extrapolation to the hard X-ray power-law, and which we speculate may be similar to the continuum component observed in BL Lac objects in their quiescent periods. From the \\sax\\ and \\asca\\ data, we find that the slope of the hard X-ray power-law depends very much on the form of the soft component which is assumed. For our best-fitting models, it lies somewhere between $\\alpha$=0.7 and 1.3 and thus may not be significantly softer than AGN in general. ",
        "introduction": "The narrow-line Seyfert 1 (NLS1) type of AGN has proven to be a valuable and a fascinating resource for the study of the class as a whole. They have very strong soft X-ray excesses \\citep{puc92,bol96}, are often highly variable \\citep{bra96,bol97} and, in some cases, the optical/UV big blue bump (BBB) component is so hot that it is shifted into the UV/EUV regime, leaving a bare, power-law-like continuum component in the optical \\citep{puc95}.  One of the earliest hypotheses put forward to explain the strong and variable ultra-soft X-ray excesses, was a high mass accretion rate onto a relatively low-mass black hole \\citep{pou95}. This was originally proposed for \\rej, by analogy with the properties of Galactic Black Hole Candidates (GBHCs). An alternative model was that such systems were geometrically-thick accretion disks (ADs) viewed face-on, lying co-planar with a flattened broad line region [BLR; \\citet{puc92}].  In our initial study of \\rej\\ \\citep{puc95}, we showed that the overall IR to X-ray spectrum compared well with the combination of a geometrically-thin, optically-thick AD and an underlying power-law with a slope, $\\alpha$=1.3 ($\\alpha$ is defined throughout such that $F_\\nu\\propto\\nu^{-\\alpha}$). The Kerr black hole mass was 7$\\times10^5$ M$_\\odot$, the accretion rate was 0.073 M$_\\odot$ yr$^{-1}$ and the disk was viewed at an angle of 60$^\\circ$ from its axis. However, this was for illustrative purposes only and no constraints were placed on these parameters. Thus, while this might have suggested the presence of a low-mass black hole for \\rej, consistent with the GBHC analogy, solutions at higher black hole mass could not be ruled out.  With the launch of {\\sl Beppo-SAX} and the inclusion of \\rej\\ in the NLS1 core program, came the opportunity to place meaningful constraints on black hole mass (M$_{bh}$), accretion rate (\\.M) and inclination. \\rej\\ was also due to be observed with {\\sl HST} within a few months of {\\sl Beppo-SAX}, and an optical spectrum was also scheduled in {\\sl WHT} service time within a few weeks. Thus a quasi-simultaneous spectrum, providing the most complete coverage possible from the optical to mid X-rays was available. Simultaneity is important for a NLS1 due to the variable nature of the class, although \\rej, contrary to its counterparts, has shown remarkable stability. Thus these data have provided the best opportunity yet to fit AD models and provide constraints on the defining parameters of an NLS1. Using a combination of a power-law and an optically-thick, geometrically-thin Kerr AD model, we present fits to the observed spectrum of \\rej, measuring M$_{bh}$ and \\.M and constraining the inclination for the first time, and discuss the results in the context of the GBHC model.  The improved UV and optical data presented here have also allowed us to re-visit the issue of what is producing the optical/UV continuum in this object. Our previous work had suggested a power-law-like continuum rising towards the red, which was not due to the host galaxy or the BBB. We have separated out the galaxy component from the spatially-resolved optical spectrum obtained and, when combined with the {\\sl HST} data, examined the pure nuclear component with a greater degree of clarity. The results and their implications for the production of the optical continuum in AGN are also discussed in this paper. ",
        "conclusions": "We have presented a comprehensive study of the narrow-line Seyfert 1 galaxy \\rej, spanning the optical to X-rays and six years of observations. Focusing on quasi-simultaneous \\sax, \\hst, \\euve\\ and optical data taken in early 1997, we have {\\sl (1)} measured the soft and hard X-ray components simultaneously for the first time using \\sax; {\\sl (2)} placed tight constraints on the form of the big blue bump using \\hst, \\euve\\ and \\sax; {\\sl (3)} separated out the nuclear component from the host galaxy in the optical; and {\\sl (4)} fitted the optical to X-ray nuclear continuum with the combination of an accretion disk and an underlying power-law, to constrain the black hole mass, accretion rate and inclination of the accretion disk in \\rej.  Coverage of the widest possible range in X-rays, particularly at soft energies, is of prime importance for NLS1s, because of their strong ultra-soft components. These \\sax\\ data, which cover the 0.1 to 10~\\keV\\ range, provide significant constraints on the shape, and thus the physical nature, of the X-ray spectrum in \\rej.  Our analysis of the \\sax\\ spectrum reveals a significant hardening of the X-ray spectrum above $\\sim$3~\\keV\\ and that two blackbodies, with $kT_{\\rm eff}$=60~\\eV\\ and 160~\\eV\\ best represent the ultra-soft X-ray component, consistent with multi-temperature emission from an optically-thick accretion disc. We do not find any compelling evidence to suggest that the hard X-ray power-law in \\rej\\ is significantly softer than that of non-NLS1 AGN. There is the suggestion of a low-energy flattening in the \\sax\\ data, which was also observed in the earlier \\rosat\\ PSPC spectrum. This in turn implies a high temperature for the inner edge of the accretion disk, and thus a low black hole mass (which is also borne out by the modeling).  The results of fitting the optical to X-ray nuclear continuum have shown that the data prefer relatively high mass accretion rates ($\\sim$0.3 to 0.7 $\\dot M_{\\rm Edd}$) onto a low mass black hole (M$_{bh}\\sim$2-10$\\times 10^6$ M$_{\\odot}$) at high inclination angles (ie. preferably edge-on, $\\sim 60-70^\\circ$ away from the axis of the disc). The very low intrinsic columns implied by the X-ray fits suggest that any molecular torus must either be optically thin, geometrically thin, or lie out of the plane of the accretion disk so that it does not obscure our line of sight.  The optical to X-ray nuclear continuum {\\sl requires} the presence of a flat, power-law-like component in the optical/UV. This is consistent with an extrapolation to the hard X-ray spectrum, although an independent origin for the hard X-ray power-law cannot be ruled out. The existence of such a component has been proposed before, but we believe that this is the first direct observation, by virtue of the unusually hot big blue bump in \\rej. It appears to be stable and unpolarized and we speculate that it may be a BL~Lac-type component which is currently in the quiescent stage of its duty-cycle."
    },
    "0011/astro-ph0011340_arXiv.txt": {
        "abstract": "We present contemporary infrared and optical spectroscopic observations of the type~IIn SN~1998S covering the period between 3 and 127 days after discovery. During the first week the spectra are characterised by prominent broad H, He and C~III/N~III emission lines with narrow peaks superimposed on a very blue continuum (T$\\sim$24000~K). In the following two weeks the C~III/N~III emission vanished, together with the broad emission components of the H and He lines. Broad, blueshifted absorption components appeared in the spectra.  The temperature of the continuum also dropped to $\\sim$14000~K. By the end of the first month the spectrum comprised broad blueshifted absorptions in H, He, Si~II, Fe~II and Sc~II. By day~44, broad emission components in H and He reappeared in the spectra. These persisted to as late as $\\sim$100-130 days, becoming increasingly asymmetric.  We agree with Leonard {\\it et al.} (2000) that the broad emission lines indicate interaction between the ejecta and circumstellar material (CSM) emitted by the progenitor.  We also agree that the progenitor of SN 1998S appears to have gone through at least two phases of mass loss, giving rise to two CSM zones. Examination of the spectra indicates that the inner zone extended to $\\leq$ 90~AU, while the outer CSM extended from 185~AU to over 1800~AU.  We also present high resolution spectra obtained at 17 and 36 days. These spectra exhibit narrow P~Cygni H~I and He~I lines superimposed on shallower, broader absorption components. Narrow lines of [N~II], [O~III], [Ne~III] and [Fe~III] are also seen.  We attribute the narrow lines to recombination and heating following ionisation of the outer CSM shell by the UV/X-ray flash at shock breakout.  Using these lines we show that the outer CSM had a velocity of 40--50~km/s. Assuming a constant velocity, we can infer that the outer CSM wind commenced more than 170 years ago, and ceased about 20~years ago, while the inner CSM wind may have commenced less than 9 years ago. During the era of the outer CSM wind the outflow from the progenitor was high - at least $\\sim2\\times 10^{-5}$M$_{\\odot}$~yr$^{-1}$.  This corresponds to a mass loss of at least $\\sim$0.003~$M_{\\odot}$, suggesting a massive progenitor. The shallower, broader absorption is of width $\\sim$350~km/s and may have arisen from a component of the outer CSM shell produced when the progenitor was going through a later blue supergiant phase. Alternatively, it may have been produced by the acceleration of the outer CSM by the radiation pressure of the UV precursor.  We also describe and model first overtone emission in carbon monoxide observed in SN~1998S.  We deduce a CO mass of $\\sim$ 10$^{-3}$ M$_{\\odot}$ moving at $\\sim$2200~km/s, and infer a mixed metal/He core of about 4~$M_{\\odot}$, again indicating a massive progenitor.  Only three core-collapse supernovae have been observed in the $K$-band at post-100~days and all three have exhibited emission from CO. ",
        "introduction": "Core-collapse supernovae (SNe) are believed to arise from massive progenitors (M$\\geq$8-10M$_\\odot$).  The lower-mass ($\\sim$ 8-10 M$_{\\odot}$) and higher-mass ($>20$ M$_{\\odot}$) progenitors experience heavy mass loss during the final stages of their evolution, several solar masses being ejected via a range of mass-loss rates. Consequently, in the vicinity of some core-collapse supernovae, dense circumstellar material (CSM) would be distributed according to the mass-loss history of the progenitor. The subsequent interaction of the freely-expanding supernova ejecta with the slowly-moving CSM generates a fast shock wave in the CSM and a reverse shock wave in the ejecta. The shocked regions emit high-energy radiation.  The intensity of this emission depends on the density of the CSM and the ejecta, and the shock acceleration of the ejecta during the explosion (Chevalier \\& Fransson 1994). If the density of the CSM is small then the effects of the interaction only become significant several years after the explosion when the supernova has faded. However, if the CSM near the supernova is relatively dense, strong CSM-ejecta interaction can begin shortly after the explosion.  This is the case for type IIn (n=narrow line) supernovae. These events exhibit narrow emission lines in their spectra superimposed on broader emission profiles.  They also exhibit a strong blue continuum (Schlegel 1990). However, the broad P~Cygni absorption components typical of normal type~II SNe are weak or absent in type~IIn SNe.  The presence of variable, narrow line emission is a direct manifestation of the excitation of the dense CSM by the SN radiation. In addition, the presence of broad H$\\alpha$ emission without a broad P Cygni absorption component is a clear indication that the observed broad line emission is powered by the ejecta-wind interaction (Chugai 1990).  Consequently, type~IIn supernovae can provide unique information about the progenitor and its later evolution, through the observed properties of their CSM. In addition, the study of the interaction of these supernovae with the CSM can provide vital clues about galaxy evolution and the nature of active galactic nuclei ({\\it c.f.}  Terlevich {\\it et al.}  1992).  SN~1998S is the brightest type~IIn event ever observed.  It was discovered on 1998 March 2.68 UT in the highly-inclined Sc galaxy NGC 3877 by Z. Wan (Li \\& Wan 1998) at a broadband (unfiltered) optical magnitude of +15.2.  The supernova is located at 16$^{\\prime\\prime}$ west and 46$^{\\prime\\prime}$ south of the nucleus.  By March 18.4 it had brightened to $V=+12.2$ (Fassia {\\it et al.} 2000). In a prediscovery frame obtained on 1998 February 23.7 there was no evidence of the supernova, down to a limiting apparent magnitude of $\\sim$+18 (Leonard {\\it et al.} 2000, IAUC 6835).  We can thus assume that SN~1998S was discovered within a few days of the shock breakout. In the present paper we have adopted the discovery date 1998 March 2.68 UT = JD 2450875.2 as epoch 0 days, t$_{0}$, and express  all other epochs relative to this fiducial date.  Optical spectra obtained about day~3 (March~5--6) by Filippenko \\& Moran (1998), Huchra (Garnavich {\\it et al.} 1998) and ourselves (see $+$3.3 d spectrum in Figure~\\ref{figop}) showed prominent H and He emission lines with narrow peaks and broad wings superimposed on a blue continuum.  As mentioned above these narrow lines indicate the presence of a dense CSM in the vicinity of the supernova.  Using optical spectropolarimetry obtained at 5~days, Leonard {\\it et al.} (2000) [L00] deduced that the CSM is asymmetrically distributed. Bowen {\\it et al.} (2000) have presented high resolution spectroscopy of interstellar and circumstellar lines towards SN~1998S.  They suggest that the CSM comprises a dense shell expanding at ~50~km/s with a more highly ionised shell moving at $\\sim$300~km/s.  They also estimate that SN~1998S is at a real distance of $\\sim$10~kpc from the nucleus and deduce from the interstellar lines that it lies on the far side of the galaxy disk.  Gerardy {\\it et al.}  (2000) [G00] presented near-IR spectra of SN~1998S spanning 95--355~days post-maximum light. They identified emission from carbon monoxide. They also suggested that late-time multi-peak H and He line profiles in their optical and IR spectra indicate emission from a disk-shaped or ring-shaped circumstellar component.  In addition, their t$\\geq$225~d near-infrared spectra exhibit a continuum that rises towards the longer wavelengths. They propose that the rising continuum is likely due to dust heated by the interaction of the ejecta with the CSM.  In Fassia {\\it et al.} (2000) we presented contemporaneous optical and infrared photometric observations of SN~1998S covering the period between 11 and 146 days after discovery.  Using the interstellar Na~I~D lines we derived an extinction $A_{V}=0.68^{+0.34}_{-0.25}$ mag.  We also examined the evolution of the total luminosity and found that during the first month the luminosity decreased very rapidly {\\it viz.} $\\sim$0.08~mag/day.  Subsequently (30-70 days) the decline rate decreased, resembling that of type~IIL supernovae {\\it viz.} 0.05 mag/day. By day~$\\sim$100 the decline rate slowed to 0.01 mag/day matching the radioactive decay of $^{56}$Co.  From the bolometric luminosity after $\\sim$100 we estimate that 0.15$\\pm$0.05~M$_{\\odot}$ of $^{56}$Ni were produced in the explosion.  Furthermore, we discovered that as early as day 130 the supernova exhibited an astonishingly high infrared (IR) excess, $K-L'=+2.5$.  We argue that this excess is due to dust grains in the vicinity of the supernova. However, the physical extent of this early IR luminosity source was so large that the emission must have come from pre-existing dust in the CSM, possibly powered by X-rays from the ejecta-CSM interaction.  In this paper we present optical and infrared spectroscopy of SN~1998S covering the period 3--127~days after discovery.  The observations and the reduction procedure are presented in section~2.  In section~3 we discuss possible line identifications. In section~4 we present an overview of the spectral behaviour of SN1998S and derive constraints about the nature and the characteristics of the CSM. We also analyse and model the first overtone of the CO emission. The work is summarised in section~5. ",
        "conclusions": "We have presented and discussed optical and infrared spectra of the type~IIn SN~1998S, covering epochs from a few days after explosion to over 100~days later.  Our observational coverage of SN~1998S makes a significant contribution to the study of the type~IIn phenomenon.  Of particular note is our acquisition of (a) contemporary spectra in both the optical and IR bands at a range of epochs, and (b) the most extensive set of high resolution spectra ever for this type of supernova event.  The spectroscopic evolution of SN~1998S was complex. It can be understood in terms of the interaction of the supernova with a two-component progenitor wind which we have referred to as the ICSM and OCSM.  Collision of the ejecta with the ICSM accounts for the early spectral features. From the fading of the broad emission components by $\\sim$14d we deduce that the outer boundary of this wind lay at less than 90~AU from the centre.  Estimation of the inner wind velocity is difficult due to (a) the ongoing interaction of the shock with this wind even at the earliest of times, and (b) the lack of high resolution spectra at these times. We can only say that the early-time spectra exhibited features of width less than 400~km/s.  However, these may have arisen in the OCSM.  If, like L00, we assume that the inner and outer winds had similar velocities then from our high resolution measurements of the OCSM we can infer that the ICSM wind commenced less than 9~years ago. However, the possibility of an RSG-BSG evolution would argue against this assumption.  Examination of the spectra indicates that the OCSM extended from 185~AU to over 1800~AU.  Our high resolution spectra have revealed that the OCSM has a velocity of 40--50~km/s.  Assuming a constant velocity during this time, we can infer that it commenced more than 170 years ago, and ceased about 20~years ago.  During this period the outflow was high - at least $2 \\times 10^{-5}$M$_{\\odot}$~yr$^{-1}$, corresponding to a mass loss of at least 0.003~$M_{\\odot}$.  An outflow of this strength and velocity is similar to those seen in cool supergiants.  The broader absorption feature ($\\sim$350~km/s) in H and He~I may have arisen from a component of the outer CSM shell produced when the progenitor was going through a later blue supergiant phase. Alternatively, it may have been due to UV-precursor radiative acceleration of the inner part of the OCSM.  Our analysis of the CO emission together with the bolometric light curve Fassia {\\it et al.} 2000) also indicates a massive progenitor, with a mixed metal/He core of $M\\sim4~M_{\\odot}$ i.e. comparable to that of SN~1987A.  SN~1998S is only the third core-collapse SN for which post-100~day $K$-band spectra have been published, and yet all SN three events have exhibited first-overtone CO emission. (Fassia {\\it et al.} (1998) also published IR spectra for the core-collapse SN~1995V.  No CO was identified but the latest $K$-band spectrum was only at 69~days.)  A picture is therefore beginning to emerge where CO plays a ubiquitous role in the evolution of core-collapse SNe at late times.  In particular, the powerful cooling property of CO is likely to lead to conditions in which dust may condense in the ejecta. Excess IR emission has been reported in SN~1998S on days~225, 260 and 255 (G00), day~253 (Garnavich {\\it et al.} 1998) and for a series of dates to as late as day~691 (Fassia et al. 2000), and this may be due to dust condensation.  (An IR excess at 130~days was also reported by Fassia {\\it et al.} (2000) but they argue that it was unlikely to have arisen from dust condensation).  In a future paper (Fassia {\\it et al.} in preparation) we shall examine in detail the evolution of the late-time IR excess.  L00 used spectropolarimetry at early times to infer significant asymmetry in the OCSM and in the ejecta-ICSM interface.  While our discussion of the high resolution OCSM lines assumed spherical symmetry, we have not provided quantitative explanations for the relative shifts between the allowed and forbidden lines.  It may be that this will also require the introduction of an asymmetric model. G00 also deduced asymmetry in the OCSM, on the basis of the broad line profiles of H and He on days~240, 275, 370.  We shall discuss these profiles in a future paper."
    },
    "0011/astro-ph0011495_arXiv.txt": {
        "abstract": "We discuss the different definitions of the mass of a halo in common use and how one may convert between them.  Using N-body simulations we show that mass estimates based on spherical averages are much more tightly correlated with each other than with masses based on the number of particles in a halo. The mass functions pertaining to some different mass definitions are estimated and compared to the `universal form' of Jenkins et al.~(\\cite{JFWCCEY}). Using a different simulation pipeline and a different cosmological model we show that the mass function is well fit by the Jenkins et al.~(\\cite{JFWCCEY}) fitting function, strengthening the claim to universality made by those authors. We show that care must be taken to match the definitions of mass when using large N-body simulations to bootstrap scaling relations from smaller hydrodynamical runs to avoid observationally significant bias in the predictions for abundances of objects. ",
        "introduction": "One of the most fundamental predictions of a theory of structure formation is the number density of objects of a given mass, the mass function, at a given redshift.  Accurate mass functions are used in a number of areas in cosmology; in studies of galaxy formation, in measures of volumes (e.g.~galaxy lensing) and in attempts to infer the normalization of the power spectrum and the density parameter from the abundance of rich clusters. In the latter case the mass function is the point of contact allowing us to bootstrap the excellent statistics of large N-body simulations with observable properties of clusters normalized for example by hydrodynamic simulations of smaller volumes.  In this way we can obtain reliable estimates of e.g.~the number of clusters as a function of temperature and redshift, which can in turn be used to constrain the matter density, the normalization of the power spectrum and the statistics of the initial density field.  Several different definitions for the ``mass'' of a halo are in common use, each having different advantages. For example there are at least 3 different definitions of mass for the mass-temperature relation as computed from hydrodynamical simulations of galaxy clusters and all are different from the definition of mass commonly used in the mass function which is itself different from the mass usually employed in analytic studies based on the Press-Schechter~(\\cite{PS}; hereafter PS) theory. Observational data have improved to the point where it is important to distinguish between these different definitions of mass, lest we bias our theoretical predictions.  We give an example of how such a conversion can be made, at least approximately, for a certain class of mass estimator. ",
        "conclusions": "We have shown that different definitions of the ``mass'' of a halo exist, and have different strengths and weaknesses.  It is important to be consistent when combining relations which use different definitions of mass, and we have given an approximate method for converting between some commonly used mass estimators. Mass estimates based on spherical averages are much more tightly correlated with each other than with the mass obtained simply by summing the particles in the group, and can be quite well estimated by assuming a `universal' spherical profile (\\S\\ref{sec:corr}). The hydrodynamical simulations which calibrate observables as a function of cluster mass typically use such a spherically averaged mass definition.  Unfortunately these definitions are not very tightly correlated with the particle based mass used in the universal mass function of halos reported by Jenkins et al.~(\\cite{JFWCCEY}).  We have argued that this is because, with a linking length of $b=0.2$, FOF is merging neighbouring halos.  Such a problem would be mitigated by reducing $b$, but it has not been demonstrated that the mass function is universal for $b\\ne 0.2$.  In the cosmological model we have simulated, the fitting form of Jenkins et al.~(\\cite{JFWCCEY}) provides a good match to the mass function of our halos, strengthening the claim of those authors that it is universal. However we find the best fit when the mass estimator used is the top-hat virial mass, rather than the FOF mass. While we have not investigated other cosmological models, we expect that the profiles (from the point of view of estimating masses) will not be very cosmology dependent. It would be interesting to see whether the Jenkins et al.~(\\cite{JFWCCEY}) mass function is `universal' if one uses the (cosmology dependent) top-hat virial mass of the halos.  \\bigskip I would like to thank V.~Springel for the use of his FOF group finder, J.~Mohr for useful conversations on this issue, and C.~Metzler for encouraging me to write it up. M.~White was supported by the US National Science Foundation and a Sloan Fellowship. Parts of this work were done on the Origin2000 system at the National Center for Supercomputing Applications, University of Illinois, Urbana-Champaign."
    },
    "0011/astro-ph0011176_arXiv.txt": {
        "abstract": "\\vskip .2 truein Distance--redshift relations are given in terms of associated Legendre functions for partially filled beam observations in spatially flat Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) cosmologies. These models are dynamically pressure-free, flat FLRW on large scales but, due to mass inhomogeneities, differ in their optical properties. The partially filled beam area-redshift equation is  a Lame$^{\\prime}$ equation for arbitrary FLRW and is shown to simplify to the associated Legendre equation for the spatially flat, \\ie $\\OO=1$ case. We fit these new analytic Hubble curves to recent supernovae (SNe) data in an attempt to determine both the mass parameter $\\OM$ and the beam filling parameter $\\nu$. We find that current data are inadequate to limit $\\nu$. However, we are able to estimate what limits are possible when the number of observed SNe is increased by factor of 10 or 100, sample sizes achievable in the near future with the proposed SuperNova Acceleration Probe satellite. ",
        "introduction": "\\label{sec-intro} Distance-redshift or equivalently the Hubble curve is critical in determining current values of the cosmological parameters $H_0, \\OM$, and $\\OL$. Conversely, current values of these three parameters determine the large scale dynamics of the Universe into the distant past. A complication occurs when attempting to determine these parameters from high $z$ comparisons to the standard Hubble curve.  The standard Hubble curve is a theoretical quantity computed assuming all gravitating matter is homogeneously distributed; whereas, observational data is taken in the real inhomogeneous Universe. In an inhomogeneous universe an observing light beam is lensed by inhomogeneities located external to, but near the light beam, and defocused (relative to the standard Hubble curve) by the less than average matter density within the beam. The simplest way to take into account these effects is to correct all beams for the missing homogeneous matter but correct for lensing only when necessary. This procedure requires the introduction of one additional parameter, \\eg a filling parameter $\\nu, \\ 0\\le\\nu\\le2$ defined by the fraction of inhomogeneous matter $\\rho_I/\\rho_0 \\equiv\\nu(\\nu+1)/6 \\le 1$ excluded from observing beams ($\\nu=0$ is the standard 100\\% filled beam FLRW case and  $\\nu=2$ is the empty beam case). When observing high $z$ objects ($z\\sim 1$) the reader can think of the parameter $\\nu$  as representing matter that exists in galaxies but not in the intergalactic medium. To find the theoretical Hubble curve for observations in such a universe one must solve the geometrical optics equation [see \\cite{KR98}] given as equation (\\ref{Area}) in the next section. This equation is actually equivalent to the Lame$^{\\prime}$ equation for general FLRW but as pointed out by \\cite{KKT} reduces to the associated Legendre equation (\\ref{Legendre}) for the special case considered here, $\\OO=1$. In \\S\\,\\ref{sec-lumdist} we solve this equation using appropriate boundary conditions and give the Hubble curve in terms of associated Legendre functions (eq.[\\ref{Pans}]) as well as in terms of hypergeometric functions (eq.[\\ref{2F1ans}]). In \\S\\,\\ref{sec-fit} we fit this new Hubble curve to data for 60 supernovae (SNe) from the Supernova Cosmological Project (SCP) and from the Cala$^{\\prime}$n/Tololo Supernova Survey (CTSS) in an attempt to determine the mass parameter $\\OM$ and the filling parameter $\\nu$. In \\S\\,\\ref{sec-conclusions} we give some concluding remarks. ",
        "conclusions": "\\label{sec-conclusions}  We have given useful forms for the luminosity distance in the currently relevant inhomogeneous $\\OO=1$ FLRW cosmologies.  These cosmologies are all dynamically FLRW in the large but differ in how gravitating matter affects optical observations.  A beam filling parameter $\\nu, \\ 0\\le\\nu\\le2$ allows the matter to vary from completely transparent and homogeneous to completely inhomogeneous and exterior to any observing beam. In order to determine the values of the cosmic parameters ($\\OM$, $\\OL$) from Hubble curves at high redshift, the value of $\\nu$ must also be constrained.  For fixed values of ($\\OM, \\OL$), increasing $\\nu$ from 0 (totally homogeneous universe) to 2 (totally clumped) increases the distance moduli of points on the Hubble curve, {\\it especially at higher redshifts} as pointed out in \\S \\ref{sec-fit}.  When observational error is taken into account, the problem of using standard candles at high redshift while ignoring $\\nu$ to obtain $\\OL$ will become particularly confounding.  The current sample of SNe Ia at $z < 1$ fail to constrain the value of the beam-filling parameter $\\nu$.  Samples 10 to 100 times larger than the current sample, and in the same redshift range, will constrain $\\nu$.  In order to unambiguously determined ($\\OM$,$\\OL$) from even higher redshift observations like those planned in the future, the distance-enhancing effect of $\\nu$ must be accounted for in the luminosity distance formulae."
    },
    "0011/astro-ph0011426_arXiv.txt": {
        "abstract": "We discuss the systematic effects arising from the {\\it cosmological redshift-space (geometric) distortion} on the statistical analysis of isodensity contour using high-redshift catalogs. Especially, we present a simple theoretical model for isodensity statistics in cosmological redshift-space, as a generalization of the work by Matsubara (1996). The statistical quantities considered here are the two- and three-dimensional genus of isodensity contour, the area of surface, the length of contour intersecting with a plane and the number of the crossing points of isodensity contour on a line. We give useful analytic formulae for the isodensity statistics, which take into account the corrections from the geometric distortion, the nonlinear clustering and the nonlinear velocity distortion phenomenologically. We then demonstrate how the geometric distortion and the nonlinear corrections alter shapes of the statistical quantities on the basis of plausible cosmological models. Our results show that the nonlinear correction can be sensitive to a choice of the redshift-space coordinate as increasing the redshift. The low-dimensional quantities such as two-dimensional genus systematically yield anisotropy due to the geometric and velocity distortions and their angle-dependence shows the $10\\sim20\\%$ difference of amplitude. Sensitivity for typical high-redshift samples are estimated in an analytic manner, and the influence of the light-cone effect for the isodensity statistics is also discussed. A simple estimation suggests that the systematic effects of geometric and redshift-space distortions can become comparable or could be dominant to the statistical error of deep cluster samples and future high-redshift galaxy surveys. These systematic effects might be a useful tool in probing the cosmological model of our universe. ",
        "introduction": "\\label{sec: intro} Recent observations of high-redshift objects have revealed dynamical aspects of the evolving universe. Measurements of Lyman-break galaxies provide evidences for the early galaxy assembly at the epoch $z\\sim 3$, which put a serious constraint on the structure formation scenario (Steidal et al. 1998; Giavalisco et al. 1998). Clusters up to $z\\sim 1$ have been observed, and the abundance suggests the universe with a low density parameter $\\Omega \\sim 0.3$ (e.g, Carlberg et al. 1997). Moreover, the upcoming wide-field surveys such as the Sloan Digital Sky Survey (SDSS) and the Two-Degree Field (2dF) survey promise to extend the observable scale of the universe and these enormous data will precisely reveal the clustering feature of the high-redshift quasars and galaxies. Since the analysis combining different redshift data sets has a potential to break the degeneracy of the cosmological parameters, as well as to provide the strong constraints on the structure formation scenario, statistical studies using the high-redshift objects are of great importance as a complementary to the measurements of cosmic microwave background anisotropy and the detection of cosmic shear.  Recently several authors have discussed various cosmological effects arising from proper observations of the high-redshift objects (e.g, Suto et al. 1999 and references therein). Apart from the biasing of the luminous objects, the main effects are summarized as follows:  \\begin{enumerate} \\item geometric distortion : no one can obtain the correct three-dimensional map of the high-redshift objects without knowing the correct cosmological model of the universe. The uncertainty of the cosmological model may affect proper evaluation of the clustering pattern in the  {\\it cosmological redshift space} (Alcock \\& Paczy\\'nski 1979). \\item redshift-space distortion : peculiar motion of a cosmological object affects the estimation of the distance in a redshift survey, which causes the linear redshift-space distortion due to the bulk motion and the nonlinear velocity distortion due to the virialized random motion (finger-of-God effect) (Davis \\& Peebles 1983; Kaiser 1987; Suto \\& Suginohara 1991; Hamilton 1998 for a review). \\item light-cone effect : because an entire sample does not consist of objects on a constant-time hypersurface but on a light cone, the redshift evolution of the objects inevitably contaminates the data. Therefore, in a wide and deep survey, it becomes important to incorporate the light-cone effect into statistical quantities, for example, two-point correlation function and the power spectrum (Matarrese et al. 1997; Yamamoto \\& Suto 1999). \\end{enumerate}  For the popular two-point statistics, the significance of these effects have been studied in detail and the robustness of the theoretical prediction is checked by the numerical simulation (Ballinger et al. 1996; Matsubara \\& Suto 1996; Magira et al. 2000; Hamana et al. 2000). Using the theoretical formulae, the feasibility of cosmological test with geometric distortion has been investigated (Yamamoto, Nishioka \\& Taruya 2000). Further, the analysis is developed to the higher-order moments on a light-cone (Matsubara et al. 1997).   Here we extend these analyses to the isodensity statistics. That is, we investigate the influences of the above cosmological effects on the isodensity statistics, which have not been ever clarified. The isodensity statistics like the genus $G_3$ and area statistic $N_3$ have progressively become important as a quantitative measure of the topology of the large-scale structure. Analytical expressions of the isodensity statistics are derived for the Gaussian random field and some cases of non-Gaussian field in real space (e.g, Bardeen et al. 1986; Hamilton et al. 1986; Matsubara 1994, 2000; Matsubara \\& Yokoyama 1996). With these analytic formulae, the isodensity statistics computed from the redshift surveys of galaxies and/or the simulated dark matter catalogues are comprehensively understood (Colley et al. 2000 and references therein).   While most authors of these works have focused on the low-redshift objects except for the cosmic microwave background anisotropy, the analysis using the high-redshift samples is expected to provide an important clue on the origin and evolution of proto-galaxies and proto-clusters. In fact, Park et al. (2000) have recently analyzed the two-dimensional genus using the galaxies taken from the Hubble Deep Field and quantified the feature of early galaxy assembly. Remarkably, their samples are widely distributed over the redshift range $0\\simlt z \\simlt 3$, which revealed an unprecedentedly deep view of galaxy clustering. Although the analysis is still limited to the projected density field because of the large error of photometric redshift, precise measurement of redshift will enable us to quantify the various isodensity statistics defined in the three-dimensional space, as well as the quantities in the one- and two-dimensional space in near future.   Practically, the observed distribution of high-redshift objects suffers from the cosmological effects, which should be correctly taken into account when comparing such data with theoretical prediction. Then naive but important questions are as follows: how the cosmological effects affect the statistical study of isodensity contour; whether these effects are observed as signal or noise. Purpose of this paper is to address these issues by using a simple theoretical formulae for the isodensity statistics, in which the cosmological effects are faithfully incorporated. Partially, the influence of the linear (redshift-space) distortion is investigated by Matsubara (1996), though the effect turns out to be small. In his paper the corrections from the nonlinear clustering and the nonlinear velocity distortion, which become influential on small length scales, are not investigated, and neither is the geometric distortion.   In this paper, we generalize the Matsubara's results to those incorporating the geometric distortion effect and the nonlinear correction taking account of the smoothing effect. With the extended formulae, we examine how the geometric and redshift-space distortions affect the isodensity statistics qualitatively and quantitatively. We will show that the nonlinear effect can be important even for a case with large smoothing scale because the scaling effect of the geometric distortion causes apparent enhancement of the nonlinear effects, which depends on choice of the cosmological redshift space. In addition we also discuss the sampling noise and the effect of finite redshift depth for typical high-redshift samples by using our extended formulae.  We organized the paper as follows.  In section \\ref{sec: Cosmological redshift space}, the cosmological redshift-space coordinate is introduced, and the statistical quantities incorporating the nonlinear corrections are described. Then we consider the isodensity statistics in cosmological redshift space in section \\ref{sec: isodensity}. Based on the analytic formulae for the Gaussian field in appendix \\ref{appendix: gaussian results}, we present a nonlinear model of isodensity statistics. After demonstrating predictions of the model, sensitivity of the isodensity statistics and the effect of the light-cone are discussed in section \\ref{sec: discussion}. The final section \\ref{sec: conclusion} is devoted to summary and conclusion. ",
        "conclusions": "\\label{sec: conclusion} In the present paper, we have considered various observational effects on the isodensity statistics of high-redshift objects, and estimated their influences, as a generalization of the work by Matsubara (1996). We have presented a simple theoretical model incorporating the geometric distortion and the nonlinear correction arising from the gravitational growth and the finger-of-God effects, together with useful analytic formulae. From this theoretical model (\\ref{eq: G_log})--(\\ref{eq: N1_log}), we have demonstrated that the geometric distortion affects the shapes of the isodensity statistics and that the nonlinear correction can be sensitive to a choice of the redshift-space coordinate as increasing the redshift. Based on these results, we have briefly discussed the detectability of these effects by evaluating the shot-noise and the sampling noise contaminations. Furthermore, the influence of the light-cone effect has been investigated. The important conclusions are summarized: \\begin{enumerate} \\item When the cosmological redshift space $s=\\shb(z)$ (eqs.[\\ref{eq: shb}] and [\\ref{eq: z_crd}]) is chosen, the nonlinear correction has an effect on the isodensity statistics for smoothing length larger $\\rs\\simeq 10h^{-1}$Mpc even at the higher redshift $z\\simgt1$, where the isodensity statistics have asymmetric shapes. On the other hand, when the cosmological redshift space $s=\\seds(z)$ (eqs.[\\ref{eq: seds}] and [\\ref{eq: EdS_crd}]) is chosen, the nonlinear correction turns out to be small and the linear theory of redshift-space distortions using the Gaussian results (\\ref{eq: G3_Gauss}), (\\ref{eq: G2_Gauss}) and (\\ref{eq: N3_Gauss})-(\\ref{eq: N1_Gauss}) could be a good approximation for the large smoothing scales. In both cases, the geometric and velocity distortions cause the angle-dependence in the low-dimensional quantities, and the relative amplitude can be changed by $10\\sim20\\%$.  \\item The Poisson sampling error is crucial for the sparse sampling of high-redshift objects like the (SDSS/2dF) quasars. The finiteness of the sampling volume might affect the precise measurement of isodensity statistics with Lyman-break galaxies. Though more quantitative estimation of statistical errors using numerical simulations is needed, however, the isodensity statistics can be tested using clusters and high-z galaxies from future wide-field surveys. For these samples, the cosmological effects can be comparable or could be dominant to the Poisson and the sampling noise.  \\item For a sample that the survey volume is large and the cosmological objects are distributed in a wide range of redshift, the light-cone effect can be influential. However, we infer that the light-cone effect has a weak influence on the isodensity statistics as long as the time evolution of the probability distribution function is not drastic, though this picture might depend on the evolution of the biasing factor. \\end{enumerate} Because the model presented here are quite simple, the predictions should be checked by numerical simulations. The prediction can also be sophisticated by using more reliable nonlinear models deduced from the simulations. Recent development of cosmological N-body simulation enables us to obtain the enormous large volume data over a few Gpc (Colberg et al. 2000). The larger data comparable to the Hubble volume size will provide us with chances of quantitative studies for the isodensity statistics, as well as the two-point statistics, in which the light-cone effect inevitably becomes important.   Combining our theoretical template with these simulations, various cosmological implications might be yielded. One of the most interesting application is the geometric test proposed by Alcock \\& Paczy\\'nski (1979). As we have seen in Sec.\\ref{subsec: results}, the geometric distortion makes the observed structures distorted and it causes deformation of the isodensity contour. Since this effect is sensitive to the distance-redshift relation, the anisotropy of isodensity statistics might be a useful tool in investigating cosmological models, as well as the anisotropy of two-point statistics (Yamamoto, Nishioka \\& Taruya 2000). The present theoretical formulae include useful tools for quantifying the geometric distortion, as well as for testing statistical nature of biased density fields. To address these issues, however, a more quantitative analysis should be developed.  The uncertainty of the biasing can be a critical issue to clarify the topological features of large-scale structure. Throughout the analyses, we have put the linear biasing ansatz (\\ref{eq: bias}) and the evolution of the biasing simply given by the model (\\ref{eq: FryBias}) neglecting the merging and formation processes. In more realistic situation, the biasing of luminous objects is nonlinear and stochastic (Dekel \\& Lahav 1999) and the merging and formation processes might alter the evolution of biasing (e.g, Somerville et al. 2000; Taruya et al. 2000). However, recent theoretical and numerical studies suggest that the linear biasing assumption (\\ref{eq: bias}) can be a good approximation even in the quasi-linear regime (Coles, 1993; Scherrer \\& Weinberg 1998; Matsubara 1999; Taruya \\& Suto 2000; Taruya et al. 2000). Therefore, in a qualitative sense, we believe that the present analyses would have revealed the generic feature of isodensity statistics on the observed high-redshift objects. As for the evolution of biasing, the theoretical prediction based on the Press-Schechter formalism gives the biasing factor for the dark halo distribution accurately (Mo \\& White 1996; Jing 1998; Sheth \\& Tormen 1999). Incorporating this evolution model into our analytic formulae can provide a powerful tool for quantifying the distribution of clusters of galaxies. In any way the isodensity statistics may probe a distinctive aspect of the biasing factor. This will be discussed elsewhere (Hikage, Taruya \\& Suto 2000, in preparation).   \\bigskip \\bigskip  We thank Prof. Suto for useful discussion and critical comments, Prof. Matsubara for useful discussion. A.T. gratefully acknowledges support from a JSPS (Japan Society for the Promotion of Science) fellowship.  K.Y. thanks Prof. Kojima and Mr. Nishioka for useful comments. This work is supported by the Inamori foundation and in part by the Grants-in-Aid program (11640280) by the Ministry of Education, Science, Sports and Culture of Japan. \\clearpage"
    },
    "0011/astro-ph0011083_arXiv.txt": {
        "abstract": "A numerical study of nonlinear $r$-modes in isentropic, rapidly rotating relativistic stars, via 3-D general-relativistic hydrodynamical evolutions, is presented.  On dynamical timescales, we find no evidence for strong coupling of $r$-modes to other modes at amplitudes of order one or larger. Therefore, unless nonlinear saturation sets in on longer timescales, the maximum $r$-mode amplitude is of order one. An absolute upper limit on the amplitude is set by causality. Our simulations also show that $r$-modes and inertial modes in isentropic stars are discrete modes, with no evidence for the existence of a continuous part in the spectrum. ",
        "introduction": "The study of the properties of $r$-modes in rotating compact stars and their relevance to relativistic astrophysics has received considerable attention since the discovery (Andersson 1998, Friedman \\& Morsink 1998) that these modes are unstable to the emission of gravitational radiation.  The $r-$mode instability provides an explanation for the spin-down of rapidly rotating young neutron stars to Crab-like spin-periods and for the spin-distribution of millisecond pulsars and accreting neutron stars. Additionally, it is considered to be a strong source of continuous gravitational radiation (see Friedman \\& Lockitch 1999 for a review).  Moreover, if $r$-modes induce differential rotation, their interaction with the magnetic field in neutron stars can enhance the toroidal magnetic field of the star (Rezzolla, Lamb \\& Shapiro 2000).  Before the instability can have an effect on the spin evolution of a young neutron star, the $r$-mode grows to an amplitude where it is saturated by nonlinear effects. Motivated by the absence of studies of such nonlinear saturation we performed a numerical study of $r$-mode hydrodynamical evolutions in rapidly rotating relativistic stars. We tried to elucidate what is the maximum amplitude such modes can reach, before nonlinear saturation (via hydrodynamical coupling) sets in (Stergioulas \\& Font 2000). The present contribution highlights the main findings of our study. We note that the saturation is most likely to set in on a hydrodynamical timescale, although it cannot be excluded that weak hydrodynamical couplings saturate the $r$-mode amplitude on longer timescales (but shorter than the growth timescale due to gravitational radiation reaction). However, at present, those long timescales cannot be achieved accurately in 3-D simulations, even with the largest available supercomputers. ",
        "conclusions": "\\begin{figure} \\centerline{ \\psfig{file=font_f1.eps,width=6.2cm,height=6.4cm} \\psfig{file=font_f2.eps,width=6.2cm,height=6.85cm} } \\caption{{\\it Left}: Evolution of the axial velocity in the equatorial plane for an amplitude of $\\alpha=1.0$, at $r=0.75R$. The evolution is a superposition of (mainly) the $l=m=2$ $r$-mode and several inertial modes. The amplitude of the oscillation decreases due to numerical (finite-differencing) viscosity of the code. {\\it Right}: Fourier transform of the velocity time-evolution, showing the frequencies of the modes in the inertial frame. The frequencies are the same at different radii, which implies a discrete spectrum.} \\label{fig_2} \\end{figure}  Figure \\ref{fig_2} (left panel) displays the evolution of the axial velocity in the equatorial plane ($v^z$ along the $y$-axis) at a radius of $r\\sim 0.75R$.  The evolution is a superposition of several modes, the $l=m=2$ $r$-mode being the dominant component. The chosen amplitude of the eigenfunction is $\\alpha=1.0$ (in the Newtonian limit, our definition of the amplitude agrees with that of Owen et al 1998). The perturbed star is evolved for more than 25ms (26 $l=m=2$ $r$-mode periods), during which the amplitude of the oscillation decreases due to numerical viscosity only. With much larger amplitudes the evolution is still similar to that in Figure \\ref{fig_2}, with no sign of nonlinear saturation of the $r$-mode amplitude on a dynamical timescale. Therefore, unless nonlinear saturation sets in at timescales much longer than the dynamical one, unstable $r$-modes could be driven to large amplitudes (of order one) by gravitational radiation reaction.  An absolute upper limit on the $r$-mode amplitude is set by causality, requiring $\\sqrt{v_iv^i}<c$.  A Fourier transform of the time-evolution, as a function of the frequency in the inertial frame, is shown in the right panel of Figure \\ref{fig_2}. It reveals that the initial data we are using, excite mainly the $l=m=2$ $r$-mode ($r_2$), with a frequency of 1.03 kHz and, with smaller amplitudes, several inertial modes ($i_3$, $i_4$, $i_5$) and higher harmonics. The frequencies of the various modes are the same at any given point inside the star, which indicates that the spectrum is a sum of discrete modes, in agreement with the conclusions of Lockitch, Andersson \\& Friedman (2000).  We have also investigated the possible appearance in our evolutions of the kinematical drift reported by Rezolla et al (2000). In a rotating star with a poloidal magnetic field, this kinematical drift may wind up the magnetic field lines and limit the effect of the $r$-mode instability. Our results do indicate that the perturbed star is rotating slower near the surface (compared to the unperturbed star) and that this drift scales roughly as $\\alpha^2$ for amplitudes of order $\\alpha \\sim 1.0$, although its magnitude is significantly smaller than that estimated by Rezzolla et al (2000).  Our finding that gravitational radiation could drive unstable $r$-modes to a large amplitude implies that $r$-modes can easily melt the crust in newly-born neutron stars (Lindblom, Owen \\& Ushomirsky 1999), leaving the initial conclusions about the $r$-mode instability being a strong source of gravitational waves (see Friedman \\& Lockitch 1999) essentially unchanged."
    },
    "0011/astro-ph0011560_arXiv.txt": {
        "abstract": "Differential $E_{V-I}$ variations of up to $\\sim 0.2$ mag on a scale of arcminutes across NGC 3201 are presented in the form of an extinction map. This map, created by calculating average $E_{V-I}$ values for stars in small subregions of the field with respect to a fiducial region, greatly improves the appearance of the CMD of the cluster.  We describe how we implemented this technique in detail with our data for NGC 3201.  A comparison between our map and that of the same region extracted from the COBE/DIRBE reddening maps published by Schlegel, Finkbeiner, \\& Davis (1998)\\markcite{schlegel98} (SFD) displays larger-scale similarities between the two maps as well as smaller-scale features which show up in our map but not in the SFD map.  Several methods of determining an $E_{V-I}$ zeropoint to add to our differential extinction map are presented. Isochrone fitting proved to be the most successful one, but it produces an average $E_{V-I}$ for the cluster which is smaller than previously published values by $\\sim 1.5\\sigma$.  Finally, our results seem to support the statement by Arce \\& Goodman (1999)\\markcite{arce99} that the SFD maps overestimate the reddening in regions of high extinction. ",
        "introduction": "We are currently undertaking a survey of approximately 15 Galactic globular clusters (GCs) with the aim of identifying eclipsing binary stars around or below the turn-off-point by means of detecting brightness variations.  These systems can provide masses for Population II main sequence stars. High-quality photometry is essential for obtaining reliable values for surface temperatures, luminosities, ages, etc.  Precise extinction determinations are a critical part of obtaining these results from our photometry.  NGC 3201, by chance one of the first GCs we analyzed, is a low-latitude cluster located at $\\alpha_{2000} = 10^{h} 36^{m} 36.8^{s}$ and $\\delta_{2000} = -46^{\\circ} 24^{'} 40^{''}$ or $l = 277.2^{\\circ}$ and $b = 8.6^{\\circ}$ (Harris 1996)\\markcite{h96} with a retrograde orbit around the center of the Milky Way.  Its nearby location ($d$ = 5.2 kpc) and low concentration ($c$ = 1.31) (Harris 1996) make it an attractive target for photometric studies.  Recent such studies include Cacciari (1984),\\markcite{cacciari84} Brewer et al. (1993),\\markcite{brewer93} C\\^{o}t\\'{e} et al. (1994),\\markcite{cote94} Covino \\& Ortolani (1997)\\markcite{covino97} (hereafter C\\&O), and Gonzalez \\& Wallerstein (1998)\\markcite{gonzalez98} (hereafter GW).  Due to its low-latitude position, the effects of differential reddening across the face of NGC 3201 are quite substantial.  The existence of variable extinction was noted in practically all earlier studies of this cluster.  Cacciari (1984), for instance, finds an irregular reddening distribution with a variation of $\\Delta E_{B-V}$ = 0.03 mag in addition to a mean value of $E_{B-V}$ = 0.21. GW report a range in $E_{B-V}$ of as much as 0.1 mag across NGC 3201.  Using spectroscopic data for 18 giants in NGC 3201, GW correct for the color spread by modelling the reddening in $E_{B-V}$ as a plane in $\\alpha$ and $\\delta$.  Recently published dust infrared emission maps by Schlegel, Finkbeiner, \\& Davis (1998)\\markcite{schlegel98} (SFD hereafter) seem to indicate, however, that the dust distribution in the region of NGC 3201 is too clumpy to be fitted by a linear function in position, even on the scale of arcminutes (see Fig. ~\\ref{schlegelmap}).  In this work, we attempt to remove the differential reddening across the cluster by using our high-quality $VI$ photometry. Our procedure is, in principle, similar to the one used by Piotto et al. (1999).\\markcite {piotto99} It aims to find the average $\\Delta E_{V-I}$ for subregions of the cluster field of approximately 1 arcmin$^{2}$ in size with respect to a fiducial region in the cluster where little or no differential reddening is apparent and where the overall $E_{V-I}$ is small compared to the rest of the cluster.  The pixel size (resolution) of this extinction map is approximately one fourth of the size (in area) of the SFD maps.  Since the SFD maps are based on infrared emission from dust, our map will provide an independent check of whether, if at all, they tend to overestimate reddening, as was suggested by, e.g., Arce \\& Goodman (1999)\\markcite{arce99} for regions with $A_{V} > 0.5$ mag.  Our results indicate variations in $E_{V-I}$ of up to $\\sim$ 0.2 mag. Differential reddening this strong can wreak havoc with photometric and spectroscopic studies of cluster stars. The inherent uncertainties in the various parameters of the CMD (e.g., magnitude and color of the main-sequence turnoff (MSTO), horizontal branch features, etc.)  are greatly amplified. Moreover, one may ``detect'' age or surface-temperature gradients where there are none in reality. To give an example of the magnitude of the effect, using the recently published color-temperature relations by Houdashelt, Bell, \\& Sweigart (2000),\\markcite{houdashelt00} the effective temperature of a solar-metallicity, main sequence star with $V-I \\sim 0.8$ would vary by approximately 600$K$ for a differential reddening effect of $E_{V-I} \\sim 0.2$ mag, which, in turn, could lead to errors in the metallicity determination.  Using our internal dereddening technique, we obtain a high-quality, deep-photometry CMD of NGC 3201 comprised of approximately ninety 600s exposures as well as some shorter ones, all using the $V$ and $I$ bands.  Our observations and data reductions, as well as the details about our internal dereddening method, are documented in Section 2.  Section 3 contains our results concerning the reddening map and the CMD of NGC 3201.  Finally, we discuss the determination of the reddening zeropoint of our extinction map in Section 4 and give a brief summary of our work in Section 5. ",
        "conclusions": "In the process of finding eclipsing binary stars in NGC 3201, we noticed the existence of variable reddening of up to 0.2 mag in $E_{V-I}$ on a scale of arcminutes. Using our internal dereddening method outlined in Section 2.2, we created an extinction map which is shown in Figures \\ref{extmap} and \\ref{extmapgrid}. Applying the map to our raw data (Fig. \\ref{cmd_raw}) significantly improved the appearance of the CMD (Fig. \\ref{cmd_final}).  Comparison between our extinction map and the SFD map of the same region (Fig. \\ref{schlegelmap}) showed that the same larger-scale features exist in both maps. Our map displays some additional, smaller-scale features which are absent in the SFD maps (see Fig. \\ref{data-map}).  The $E_{V-I}$ zeropoint which needs to be added to the numbers in Fig. \\ref{extmapgrid} in order to get absolute $E_{V-I}$ values is 0.15. This value is below literature results (Cacciari 1984, Harris 1996) by approximately $1.5 \\sigma$, but produced by far the best VDB isochrone fit to the data.  The zeropoint determined with the help of the SFD maps gives $E_{V-I} \\sim 0.49$ as the average value across NGC 3201 which is higher than literature values which supports the statement by Arce \\& Goodman (1999) that the SFD maps overestimate the reddening in regions of high extinction.  The results from this work will be essential in our binary star research where high-quality photometry of every binary system is necessary for distance determinations. A vital condition to obtaining these measurements is of course the knowledge of the exact extinction the star under investigation is suffering.  Furthermore, studies like this may be useful in determining properties of the ISM itself such as examining a possible dependence of $R_V$ upon position in the field of view or giving insight into the distribution and properties of the dust along the line of sight."
    },
    "0011/astro-ph0011574_arXiv.txt": {
        "abstract": "First results from high-resolution coronal spectroscopy of flares with the Reflection Grating Spectrometers on board the \\textit{XMM-Newton} satellite are reviewed. Rotational modulation in the X-ray light curve of HR 1099 is discussed. Results from time-dependent spectroscopy of flares in the active stars HR 1099, AB Dor, YY Gem are also presented. Variations in the shape of the emission measure distributions, in the abundances and in the average density of the cool plasma are discussed. ",
        "introduction": "Stellar coronae often display variability in their X-ray emission. Energetic explosive events (flares), eclipses, and rotational modulation are among the most interesting features found in X-ray light curves of coronal stars. Flares are at the center of a scenario that assumes an ensemble of flares as agents of heating of stellar coronae (``microflare hypothesis''; e.g., Parker~1988). Recent results suggest that flares can contribute significantly to the coronal heating of active stars (Audard et al.~2000). In a standard model, magnetic energy is released during a flare, heating the dense layers of chromospheric material. Hot ionized material is then driven into the corona by the pressure increase and fills the coronal loops (chromospheric evaporation; see Antonucci et al.~1994).  Optically thin X-ray radiation occurs through continuum processes  and electronic transitions.  Flares are therefore important agents that bring ``fresh'' material from the chromospheric layers to the corona.  In the non-flaring Sun, a First Ionization Potential (FIP) effect is observed, with enhanced coronal abundances for elements with low FIP ($<10$~eV), while high-FIP elements show solar photospheric values. In the stellar case, evidence for the presence of the FIP effect is less clear, but has been found in some stars (e.g., Drake et al.~1995; 1997). Recently, Brinkman et al.~(2001) has found in HR 1099 evidence for an \\emph{inverse} FIP effect and a strong enhancement of Ne relative to O. Stellar coronal flares often show enhanced metal abundances (e.g., Ottmann \\& Schmitt~1996). Also, some classes of solar flares can be Ca-rich (Ca is a low-FIP element; Sylwester et al.~1984) or Ne and S-rich (Ne and S are high-FIP elements; Schmelz 1993).  In this paper, we report a summary of first results (see also G\\\"udel et al., this volume) from studies of stellar flares with the \\textit{XMM-Newton} Reflection Grating Spectrometers (RGS). Recent observations include HR 1099, AB Dor, YY Gem/Castor  (see Audard, G\\\"udel, \\& Mewe~2001 and G\\\"udel et al.~2001ab for more details). ",
        "conclusions": "Flares are frequent in the stellar coronal sources observed so far by \\textit{XMM-Newton}. Large flares occurred on HR 1099, AB Dor, and YY Gem, allowing us to perform time-dependent spectroscopy with the high-resolution RGS in the range 5--35~\\AA\\  (0.35--2.5~keV). The presence of very hot (up to several tens of MK) material has been inferred, and evidence for elemental abundance enhancements has been found. Low-FIP elements appear to increase significantly during a flare on HR 1099, while the high-FIP element Ne stays at constant abundance. This behaviour contrasts with the \\emph{inverse} FIP effect found in these coronae during quiescence. Our \\textit{XMM-Newton} observations thus indicate that flares are also important as agents that relate chromospheric/photospheric plasma with coronal plasma.  The flare EM distributions appear to be composed of a very hot plasma (up to $100$~MK) that evolves  rapidly, and of a stable quiescent plasma. This result is consistent with the absence of a change in the density of the cool material. Stellar flares are important probes for stellar coronal plasma and the coronal heating mechanism. The very high temperatures attained by these flares are unknown to solar flares. Such observations are important to extend solar knowledge to more extreme conditions appropriate for magnetically active stars."
    },
    "0011/astro-ph0011097_arXiv.txt": {
        "abstract": "We present a general linear dispersion relation which describes the coupled behavior of magnetorotational, photon bubble, and convective instabilities in weakly magnetized, differentially rotating accretion disks.  We presume the accretion disks to be geometrically thin and supported vertically by radiation pressure.  We fully incorporate the effects of a nonzero radiative diffusion length on the linear modes. In an equilibrium with purely vertical magnetic field, the vertical magnetorotational modes are completely unaffected by compressibility, stratification, and radiative diffusion.  However, in the presence of azimuthal fields, which are expected in differentially rotating flows, the growth rate of all magnetorotational modes can be reduced substantially below the orbital frequency.  This occurs if diffusion destroys radiation sound waves on the length scale of the instability, and the magnetic energy density of the azimuthal component exceeds the non-radiative thermal energy density.  While sluggish in this case, the magnetorotational instability still persists and will still tap the free energy of the differential rotation.  Photon bubble instabilities are generically present in radiation pressure dominated flows where diffusion is present.  We show that their growth rates are limited to a maximum value which is reached at short wavelengths where the modes may be viewed as unstable slow magnetosonic waves.  We also find that vertical radiation pressure destabilizes upward propagating fast waves, and that Alfv\\'en waves can be unstable.  These instabilities typically have smaller growth rates than the photon bubble/slow modes.  We discuss how all these modes behave in various regimes of interest, and speculate how they may affect the dynamics of real accretion disk flows. ",
        "introduction": "The physical state of the radiation pressure dominated, innermost regions of accretion disks around black holes has been uncertain ever since the early days of accretion disk theory.  Standard alpha disk models in which the viscous stress is assumed to scale with the radiation pressure are subject to thermal and viscous instabilities (Lightman \\& Eardley 1974, Shakura \\& Sunyaev 1976).   These instabilities are sensitive to the assumed prescription for the anomalous viscosity (e.g. Piran 1978), so it remains unclear whether or not they are actually present in real flows.  If they are, the accretion flow may adopt a radically different state from that usually envisaged in thin accretion disk theory, such as the multiphase equilibrium recently proposed by Krolik (1998).  In addition to these secular instabilities, dynamical instabilities also exist.  First and foremost, a differentially rotating flow with a negative angular velocity gradient and an initially weak magnetic field is unstable to the magnetorotational instability (MRI, Balbus \\& Hawley 1991).  The turbulence resulting from this instability is currently the most plausible mechanism known for generating the anomalous viscosity required in accretion disk models.  Dissipation of magnetohydrodynamic waves excited by this turbulence by photon diffusion and photon viscosity has recently been examined by Agol \\& Krolik (1998). Photon diffusion might also affect the linear development of the instability itself in a laminar, radiation pressure dominated flow, but this issue has not been examined previously.  Gammie (1998) has suggested that the overstable photon bubble modes discussed by Arons (1992) in the context of X-ray pulsars also exist in radiation dominated accretion flows in general.  However, his instability analysis was limited to studying a static equilibrium where the effects of differential rotation were entirely neglected. Pietrini \\& Krolik (2000) have recently investigated convective instabilities in unmagnetized, radiation pressure dominated, differentially rotating flows. They assumed a constant vertical density profile, which necessarily leads to an unstable entropy gradient.  Whether or not such unstable gradients exist in reality depends on the vertical dissipation profile, which in turn depends on the structure of the MHD turbulence.  These dynamical instabilities may all play a role at some level in the dynamics and thermodynamics of the radiation pressure dominated portion of accretion disks.  It is the purpose of this paper to provide a unified description of all three by deriving a general linear dispersion relation which incorporates them all.  Such an analysis will hopefully prove to be a useful guide to numerical simulations which explore the nonlinear development of these instabilities and their effects on the resulting turbulent state of the inner accretion disk.  This paper is organized as follows.  In section 2 we discuss our basic equations and assumptions.  Then in section 3 we focus exclusively on the MRI by deriving its dispersion relation in the absence of vertical stratification so that the other instabilities are suppressed. We are therefore able to study just the effects of photon diffusion on what is probably the most important of these instabilities.  In section 4 we then present the full dispersion relation for a vertically stratified, differentially rotating medium, and discuss its solutions. We discuss the relevance of these solutions to astrophysical accretion disks in section 5 and then summarize our conclusions in section 6. ",
        "conclusions": "We may summarize our conclusions as follows. In the absence of an azimuthal field, the vertical ($k_R=0$) MRI modes have the standard, incompressible behavior elucidated by Balbus \\& Hawley (1991).  If the azimuthal field component dominates the vertical field component and the Alfv\\'en speed is larger than the effective sound speed [cf. eq. (\\ref{vaphicond})], then the growth rate of the MRI is reduced to $\\sim\\Omega(\\cgas/\\vaphi)$, well below the usual rapid growth rate of order the orbital frequency $\\Omega$.  However, the instability still persists.  Large vertical photon fluxes in the equilibrium do not significantly affect the most rapidly growing MRI modes.  At short wavelengths where rotation is unimportant, both the fast and slow MHD waves are destabilized by sufficiently high vertical photon flux in the equilibrium.  These instabilities require that the waves propagate at some nonzero angle to the vertical, and affects upward propagating fast waves and downward propagating slow waves.  The slow waves appear to always have the larger growth rates.  At smaller wavenumbers these are well-described by Gammie's (1998) photon bubble dispersion relation (\\ref{pbdisp}), but at larger wavenumbers they asymptote to a maximum value given by equation (\\ref{slowdisp}).  At small wavenumbers where rotation is important ($k<k_{\\rm crit}$), the photon bubble growth rate is somewhat suppressed by rotation, at least at modest fluxes.  Nonzero azimuthal fields stabilize the fast wave instability and reduce the photon bubble growth rate.  In addition to the fast and slow wave instabilities, the vertical radiative flux also drives unstable modes at intermediate wavenumber which are related to Alfv\\'en waves at high wavenumber.  The ambient conditions in the radiation pressure dominated inner region of standard alpha disk models are generally in a regime where radiative diffusion can slow the growth of the MRI.  In addition, photon bubble modes are expected at short wavelengths with growth rates that can significantly exceed the orbital frequency.  It is not clear how these instabilities will evolve into the nonlinear, turbulent regime, but we speculate that the MRI will dominate the largest scales and that photon bubbles (and perhaps unstable fast and Alfv\\'en waves) may affect both the dissipation and thermal transport on smaller scales. A complete understanding of the turbulent angular momentum and heat transport will require numerical simulations.  Finally, we stress that we have concentrated here on geometrically thin accretion disk configurations where the photon flux is primarily in the vertical direction and the rotation curve is near-Keplerian. Geometrically thick or quasi-spherical, radiation pressure supported flows will have radiation pressure gradient components in the radial direction and may therefore exhibit more interesting dynamical coupling between the radiative flux and the MRI."
    },
    "0011/astro-ph0011268_arXiv.txt": {
        "abstract": "We derive the distribution of neutral and ionized gas in high redshift clouds which are optically thick to hydrogen ionizing radiation, using published data on Lyman-limit and Damped Lyman-$\\alpha$ absorption systems in the redshift range 1.75$\\le z < $3.25. We assume that the distribution of the hydrogen total (HI+HII) column density in the absorbers, $N_H$, follows a power law  $K N_{H}^{-\\alpha}$, whereas the observed HI column density distribution deviates from a pure power law as a result of ionization from a background radiation field. We use an accurate radiative transfer code for computing the rapidly varying ratio $N_{H}/N_{HI}$ as a function of $N_{H}$. Comparison of the models and observations gave excellent fits with Maximum Likelihood solutions for the exponent $\\alpha$ and for $X$, the value of log($N_{H}/N_{HI}$) when the Lyman-limit optical depth along the line of sight is $\\tau_{LL}=1$. The slope of the total gas column density distribution with its relative 3$\\sigma$ errors  is $\\alpha=2.7^{+1.0}_{-0.7}$ and $X=2.75\\pm0.35$. This value of $X$ is much lower than what would be obtained for a gaseous distribution in equilibrium under its own gravity. The ratio $\\eta_0$ of dark matter to gas density is however not well constrained since log($\\eta_0)=1.1\\pm 0.8$. An extrapolation of our derived power law distribution towards systems of lower column density, the Lyman-$\\alpha$ forest, tends to favour models with log$(\\eta_0) \\simlt 1.1$ and $\\alpha \\sim 2.7$ to 3.3. With $\\alpha$ appreciably larger than 2, Lyman-limit systems contain more gas than Damped Lyman-$\\alpha$ systems and Lyman-$\\alpha$ forest clouds even more. Estimates of the cosmological gas and dark matter density due to absorbers of different column density at $z\\sim 2.5$ are also given. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011432_arXiv.txt": {
        "abstract": "The possibility that quasi-equilibrium self-gravitating galaxy-like configurations exist in a tidal field is analyzed in this paper. More specifically, we address the question of how to predict initial configurations modeling galaxies that are able to survive environmental effects  in a dense environment for a Hubble time or so, provided thay dynamical friction is neglected. For simplicity, the configurations in the tidal field have been taken initially to be spherically symmetric and to have an isotropic velocity dispersion tensor (t-limited King spheres); they orbit inside steady state, spherical halos, as those that presumably surround compact galaxy groups and galaxy clusters. Both circular and eccentric orbits have been considered. In both cases, the initial quasi-equilibrium configurations have been built up taking into account the external tidal field produced by the halo. It modifies the escape velocity field of the configuration, compared with isolated configurations. The survival of the configurations as they orbit inside the halos has been studied through N-body simulations. As a general result, it has been found out that the bulk of the models is conserved along 12.5 Gyears of evolution, and that the  low rates of mass losses   they experience are consistent with those expected  when the adiabatic protection hypothesis is at work. So, solutions for  galaxy configurations in tidal quasi-equilibrium have been found, showing that tidal stripping in quiescent phases  does not seem to be very important, unless that  the density of the galaxy environment at its formation had been much lower than that  of the galaxy environment at the point of its orbit where the tidal perturbation is maximum. ",
        "introduction": "Dense halos of dark matter have first been detected in galaxy clusters and, then, in many galaxy groups, through observations of X-ray diffuse emission of the gas component surrounding the member galaxies (Boehringer 1997; Ponman et al. 1996). Also, some small dwarf galaxies are known to be orbiting inside the dark halos of larger ones (e.g. Mateo 1998). Interactions with such dense environments can cause these galaxies several dynamical effects (i.e. tidal heating,  mass loss from tidal stripping, energy loss from dynamical friction, among others) that could result in severe modifications, relative to isolated galaxies, of the evolutionary history of both, individual galaxies and the systems they form. In fact, some characteristic times and scales playing relevant roles in astrophysical processes, such as the typical times for decay to the center of the halos, the rates of background enrichment by processed gas and so on, could change appreciably, depending on the characteristics of the environment (e.g., Gunn \\& Gott 1972; Merritt 1985; Moore, Lake, \\& Katz 1998).  The modelization of environmental interactions has been mainly carried out through N-body simulations of the evolution of a theoretical galaxy model moving in an external field (Barnes 1985; Funato, Makino, \\& Ebisuzaki 1993; Bode, Cohn, \\& Lugger 1993; Bode et al. 1994; Garc\\'{\\i}a-G\\'omez, Athanassoula, \\& Garijo 1996; Athanassoula, Makino, \\& Bosma 1997 and references quoted therein). One of the main shortcomings of galaxy models appearing in the literature is that, in most cases, galaxies are built up as if they were {\\it isolated}. However, to properly quantify the effects of environmental interactions, it would be more convenient that the galaxy model, at the beginning of the simulation, takes into account the external forces. Otherwise, it is difficult to disentangle which effects are effectively due to interactions and which ones are spurious, due to an incorrect choice of the initial galaxy model (see G\\'omez-Flechoso \\& Dom\\'{\\i}nguez-Tenreiro 2000, hereafter GD00, for a discussion). In this paper, we will focus on the choice of the initial galaxy model  and on how this choice  affects the later evolution of the galaxy as it orbits inside a halo. The effects of the dynamical friction and the interactions between galaxies in groups and clusters will be tackled in forthcoming papers.  King models (King 1965, 1966) provide a good framework to study equilibrium self-gravitating configurations. They are based on distribution functions that depend on the  potential, $\\Phi$,  causing the forces felt by the constituent particles. So,   if the body is isolated, the potential that enters in King distribution functions is $\\Phi = \\Phi_S $, where $\\Phi_S$ is due to the mass distribution of the self-gravitating configuration (hereafter, satellite); however, when it moves through an external halo, then $\\Phi = \\Phi_S + \\Phi_{ext}$ must be used instead, where $\\Phi_{ext}$ is the potential causing the external force seen by the constituent particles. This force also determines  the limiting or tidal radius of a spherical configuration, $r_{\\rm t}$, that can be defined as the asymptotic distance at which constituent particles remain stably bound to the satellite. In spite of this, most King models found in literature are constructed on the assumption that $\\Phi = \\Phi_S $ and take $r_{\\rm t}$ as a free parameter, even if the body is not isolated (Meylan \\& Heggie 1997; but see also Heggie \\& Ramamani 1995). The tidal radius is a fundamental parameter of King models describing spherically symmetric satellites, when, as due, the effects of the external field are explicitly taken into account (hereafter, t-limited King models).  The stability of self-gravitating configurations relative to tidal perturbations is equivalent to the stability of the orbits of its constituent particles. After the pioneering work by King (1962), the problem of tidal limitations imposed on such configurations has been studied in models provided by the circular and elliptical restricted three body problems (Keenan \\& Innanen 1975; Jefferys 1976; Keenan 1981a and 1981b). When the configuration is on a circular orbit, the problem can be worked out in some detail, as the equations of motion of their constituent particles have one integral of motion, the Jacobi integral, $E_J$, that can be used to define zero velocity surfaces in the configuration space. Stability against escape for a given orbit is assumed when the corresponding zero velocity surface is closed (Spitzer 1987; note that the zero velocity surface can be open and, nevertheless, the particle does not escape over a given time interval). And so, the limiting radius, $r_{\\rm t}$, has been defined in terms of the distance, $x_e$, from the satellite center  to the inner Lagrange point of the potential $\\Phi = \\Phi_S + \\Phi_{ext}$. However, in a tidal field zero velocity surfaces are not spherically symmetric and, then, an ambiguity arises when one intends to determine the radius of a spherically symmetric body embedded in a tidal field that has not this symmetry. So, different authors make different choices. King (1962) takes $r_{\\rm t} = x_e$, while Keenan (1981b) suggests that $r_{\\rm t} = 2x_e/3$ is preferable (see also Innanen, Harris, \\& Webbink 1983; Spitzer 1987; Lee 1990; Heggie \\& Ramamani 1995); these choices correspond to  two semiaxes of the Roche surface.  No integral of motion exists when the satellite moves along a non circular orbit, as the intensity of the tidal field changes periodically with time, being maximum at pericentric passage. The question then arises of whether or not the energies of most constituent particles are  conserved to a good approximation so that  {\\it local} (i.e., depending on the satellite orbital phase) tidal  radii could be defined, as in the circular motion case, that would translate into satellites in tidal {\\it quasi}-equilibrium (note that, in the strict sense, {\\it equilibrium} configurations do not exist in a tidal field, because some degree of mass losses can never be avoided). We see that more complications are added on the theoretical side to the ambiguities appearing in the circular case. Observational data on the limiting radii of globular clusters cannot clarify the situation either, as no clear conclusion has still been  reached on their possible dependence on the cluster orbital phase (Oh \\& Lin 1992; Oh, Lin, \\& Aarseth 1992; Meziane \\& Colin 1996; Brosche, Odenkirchen, \\& Geffert 1999). In any case, results on tidal equilibrium for globular clusters, where two-body heat conduction plays an important role, could be not valid for galaxy-like configurations, that are essentially collisionless systems.  The purpose of this paper is to deepen into the understanding of tidal quasi-equilibrium for self-gravitating galaxy-like configurations, in particular to predict initial configurations for galaxy models which will survive environmental effects along a Hubble time or so. The possibility to predict self-gravitating spherical collisionless configurations in tidal quasi-equilibrium has been tested through Montecarlo realizations of the t-limited King models. They are left to evolve in the corresponding external field, which has been described by an analytical expression. Various possibilities have been explored about the matching of the internal and external field of forces at the limiting radius of the galaxy model. Most of the previous works on tidal quasi-equilibrium referred to globular clusters moving in a galactic potential. As here we are mainly concerned with the influence of dense environments (i.e., halos) on galaxy evolution, the halo density profiles used in our test are those of dark halos appearing in N-body simulations of hierarchically clustering universes (Navarro, Frenk, \\& White 1996), and the parameters of the galaxy models correspond to those of typical ellipticals.  The paper is organized as follows: in $\\S$\\ref{21} we give the general expression for the tidal field caused by a spherical static halo in the harmonic approximation. In $\\S$\\ref{proforb} we specify the models and parameters of halos, orbits and galaxies used to make our study. The results of this study are presented in $\\S$\\ref{results}. Finally, in $\\S$\\ref{sumcon}, the summary and conclusions of the work are given. ",
        "conclusions": "\\label{sumcon}  In this paper we address the issue of the existence of quasi-equilibrium self-gravitating configurations in a quiescent tidal field, that is, the possibility that self-gravitating configurations exist that are able to survive for a Hubble time or so in a given dense environment without being tidally stripped or disrupted. More specifically, we have tried to answer  the question of how to build-up such configurations with their limiting radius determined by the tidal field and their remainder characteristics described by  parameters that can be fixed a priori.  As the simplest hypotheses, the tidal field produced by a static, spherically symmetric, dense, extended halo has been considered. Also for simplicity, the configurations have been taken initially to be spherically symmetric and to have an isotropic velocity tensor (t-limited King spheres). They orbit inside the halos. Both circular and eccentric orbits have been considered. In both cases, quasi-equilibrium self-gravitating configurations have been built-up by taking as tidal radii the radii of their Roche surface along different axes, or, equivalently, by defining the escape velocity field of the configuration taking into account the requirements imposed by the tidal field produced by the external halo. The gravitational field inside the configuration is spherically symmetric, while the tidal field has no this symmetry and one has to resort to a choice of its radial component. So, an ambiguity arises when matching the internal and external fields of forces at $r_{\\rm t}$. Different possible choices have been considered. In the case of an eccentric orbit the tidal field depends on the orbital phase and a new ambiguity arises regarding the orbital position where the matching is made. Here also different possibilities have been explored.  To study the survival of the configurations, we have evolved Montecarlo realizations of t-limited King galaxy models, orbiting in the tidal field produced by a dark cluster-like halo (for circular orbits) or galaxy group-like (for eccentric orbits) halo. These galaxies have been taken to be collisionless systems, i.e., such that heat transport and relaxation cannot proceed through two-body encounters, but through the oscillations of their collective self-consistent potential.  The general result of the simulations, irrespective on how the matching has been made, is that the bulk of the models are conserved along 12.5 Gyears of evolution both for circular and eccentric orbits, even if some mass losses occur in some cases. A good linear relation between the ratio of the galaxy anomalistic period to its dynamical timescale, $T/t_{\\rm dyn}$, and the logarithm of the relative mass losses for the different galaxy models has been obtained, suggesting that  adiabatic protection is at work in these simulations. In the case of galaxies on circular orbits, the adiabatic protection hypothesis predicts zero mass losses; the results of our simulations for circular motion are  also in agreement with this prediction, once the almost imperceptible two-body effects are considered. In the case of eccentric orbits, if the galaxy configuration closely corresponds to the tidal equilibrium solution at its actual environment, no important oscillations of the potential produced by tidal forces can be expected, and, in fact, our simulations show that once the satellite particles are distributed in positions and velocities according to the equilibrium solution at one given point of the satellite orbit, they approximately remain so as the satellite orbits, except for the marginally bound particles, in which case most of them are lost to the configuration. Even if continuous mass losses are important in some cases, the system does not relax towards the equilibrium solutions corresponding to the different positions of the satellite, presumably because these different equilibrium solutions are not far enough to produce strong collective potential oscillations.  Configurations corresponding to equilibrium at pericenter, where the tidal forces are maximum, are those that suffer from less  escape (they are hyperstable solutions at the other points of the orbit) and anisotropy development. Among the simulations presented in this paper, the maximum mass losses occur when the  component of the tidal force in the ${\\bf R}_S$ direction is chosen as radial component of these forces and the initial configuration is prepared at the apocenter ($\\simeq 33.5$\\% along 12.5 Gyrs). But, as stated above, it is mainly the $T/t_{\\rm dyn}$ ratio that determines the evolution rates and mass losses.  The results described so far suggest that the configuration of spheroidal galaxies  is fixed at its formation, determined by its mass, energy content and the environment at that moment. After formation, only moderate tidally induced evolution can be expected for a galaxy living in environments  of density similar to that of its environment at formation. These results  also suggest that a continuous slow mass loss along long periods can occur, without destroying the system, if the density of the environment at formation is lower than that of the environment at galaxy pericenter passage. The galaxy will be easily destroyed, however, should it be placed at an external field whose corresponding equilibrium tidal radii is much smaller than the limiting radius of the actual galaxy (see GD00 for a discussion)."
    },
    "0011/astro-ph0011118_arXiv.txt": {
        "abstract": "beginabstract: Command not found.   Adopting new s-process nucleosynthesis scenario and branch s-process path, we calculate the heavy-element abundances of solar metallicity 3$M_{\\odot}$ thermal pulse AGB (hereafter TP-AGB) stars, and then discuss the correlation between heavy-element abundances and C/O ratio. $^{13}$C$(\\alpha,n)^{16}$O reaction is the major neutron source, which is released in radiative condition during the interpulse period, hence gives rise to an efficient s-processing that depends on the $^{13}$C profile in the $^{13}$C pocket. A second small neutron burst from the $^{22}$Ne source marginally operates during convective pulses over previously s-processed material diluted with fresh Fe seed and H-burning ashes. The calculated heavy-element abundances and C/O ratio on the surfaces of AGB stars are compared with the observations of MS, S and C (N-type) stars. The observations are characterized by a spread in neutron exposures: 0.5-2.5 times of the corresponding exposures reached in the three zones of the $^{13}$C profile showed by Fig. 1 of Gallino et al. (1998). The evolutionary sequence from M to S to C stars is explained naturally by the calculated heavy-element abundances and C/O ratio.  Then the heavy-element abundances on the surfaces of TP-AGB stars are used to calculate the heavy-element overabundances of barium stars, which are generally believed to belong to binary systems and their heavy-element overabundances are produced by the accreting material from the companions (the former TP-AGB stars and the present white dwarfs). To achieve this, firstly, the change equations of binary orbital elements are obtained by taking the angular momentum conservation in place of the tangential momentum conservation, and considering the change of $\\delta r/r$ term; then the heavy-element overabundances of barium stars are calculated, in a self-consistent manner, through wind accretion during successive pulsed mass ejection, followed by mixing. The calculated relationships of heavy-element abundances to orbital periods $P$ of barium stars can fit the observations within the error ranges. Moreover, the calculated abundances of nuclei of different atomic charge Z, corresponding to different neutron exposures of TP-AGB stars, can fit the observational heavy-element abundances of 14 barium stars in the error ranges. Our results suggest that the barium stars with longer orbital period $P>1600$ d may form through accreting part of the ejecta from the intrinsic AGB stars through stellar wind, and the mass accretion rate is in the range of 0.1-0.5 times of Bondi-Hoyle's accretion rate. Those with shorter orbital period $P<600$ d may be formed through dynamically stable late case C mass transfer or common envelope ejection. ",
        "introduction": "Since Burbidge et al. (1957) and Cameron (1957) published their creative studies, the nucleosynthesis theory has been developed in deep degree. In particular, asymptotic giant branch (hereafter AGB) stars are very important to study element nucleosynthesis and the Galactic chemical evolution because they synthesize significant parts of slow process (hereafter s-process) neutron capture elements and $^{12}$C. The products are taken out from the stellar interior, He-intershell, to the surface by the third dredge-up (hereafter TDU) process, and then are ejected into interstellar medium with the progressive stellar wind mass loss.  Our understanding of the AGB nucleosynthesis has undergone major revisions in these years. The earlier studies (Iben 1975; Truran \\& Iben 1977) illustrated that intermediate mass TP-AGB stars with $^{22}$Ne$(\\alpha,n)^{25}$Mg neutron source (the typical neutron density is least of a order of $10^{9}-10^{10} n$ cm$^{-3}$) are the suitable nucleosynthesis sites of s-process elements. But new observations shed doubts on the above idea (Busso et al. 1995 and references therein; Busso et al. 1999 and references therein). Busso et al. (1995) and Lambert et al. (1995) demonstrated that the measured abundances of Rb/Sr, the products of the branch in the s-process path at $^{85}$Kr, imply a definitely lower neutron density (typical of the order of $10^{7}n$ cm$^{-3}$), which can be provided via the reaction $^{13}$C$(\\alpha,n)^{16}$O at low temperature in the He-intershell of low mass AGB stars.  Iben \\& Renzini (1982a,b) indicated that a suitable mechanism operated in low mass stars of low metallicity to allow the formation of a semiconvective layer, hence the $^{13}$C pocket. The pocket is engulfed by the next convective pulse where $^{13}$C nuclei easily capture $\\alpha$ nuclei, release neutrons. Hollowell \\& Iben (1988) confirmed the possibility of formation of a consistent $^{13}$C pocket through a local time-dependent treatment of semiconvection. However, the semiconvection mixing mechanism was not found to work for the $^{12}$C-enriched Population I red giants, like the peculiar stars of MS, S and C stars, showing overabundances of s-process elements in their spectra.  Straniero et al. (1995) investigated the effect of a possible mixing of protons into a thin zone at the top of the carbon-rich region during each dredge-up episode, hence the formation of $^{13}$C pocket. They suggested that the $^{13}$C was completely burnt in the radiative condition, and the resulting s-process nucleosynthesis occurs during the quiescent interpulse period, instead of the convective thermal pulse. $^{22}$Ne$(\\alpha,n)^{25}$Mg was still active for a very short period during the convective pulse with minor influence on the whole nucleosynthesis. Herwig et al. (1997) and Herwig et al. (1998) supported the formation of $^{13}$C pocket via hydrodynamical calculations.  Straniero et al. (1997) adopted the above new s-process nucleosynthesis scenario to calculate the s-process nucleosynthesis of solar metallicity low mass AGB stars with $1 \\leq M/M_{\\odot} \\leq 3$, and gave the detailed results.  Recently, Gallino et al. (1998) explained further and developed the aforesaid new scenario. They divided the $^{13}$C pocket, $q$ layer, into three zones in the light of the distribution in the mass of hydrogen introduced in the $^{12}$C-rich intershell. The characteristic neutron exposures in the three layers are different. Moreover, when the nucleosynthesis occurs in a radiative layer, only the nucleosynthesis products are ingested into the convective thermal pulse, which makes the classical concept of mean neutron exposure ($\\tau_0$) become meaningless and the simple assumption of an exponential distribution of the neutron exposure fail to account for the complexity of the phenomenon (Arlandini et al. 1995, Gallino et al. 1998). Busso et al. (1999) reviewed this new s-process nucleosynthesis scenario in details.  The spectral and luminosity studies of AGB stars (including MS, S and N-type C stars) have shown that the M$\\rightarrow$S$\\rightarrow$C sequence is the result that the low-mass AGB stars have undergone carbon synthesis, s-process nucleosynthesis and the third dredge-ups (Lambert 1991). These stars are in the course of experiencing thermal pulse stage, and the original chemical abundances of their atmospheric envelope have been modified by two mixing mechanisms, namely, the first dredge-up when they became red giants and the third dredge-up when they became TP-AGB stars (Boothroyd $\\&$ Sackmann 1988a, b, c and d).  The predicted evolutionary sequence of M$\\rightarrow$S$\\rightarrow$C in the heavier$-$lighter s-element abundances relationship (here and hereafter, 'heavier' refers to the second metal peak elements: Ba, La, Ce, Nd and Sm etc.; 'lighter' refers to the first metal peak elements: Y, Zr etc.) and the heavy-element abundances$-$C/O relationship ($^{12}$C, together with the s-process elements, is dredged up from stellar interior during the third dredge-up) are important to understand the nucleosynthesis and evolution of AGB stars. Because (1) they will provide theoretical basis for the observed evolution of the sequence, (2) they can check the available theories on the evolutions of AGB stars (e.g., the beginning of the third dredge-up process, the mass and the chemical abundance of the dredge-up material, the theory of s-process nucleosynthesis, the stellar wind mass loss, and the formation of carbon stars etc.).  Busso et al. (1992) discussed the heavy-element abundances of M, MS and S stars using the thermal pulse AGB model. Busso et al. (1995) analyzed the heavy-element overabundances of carbon stars under the assumption that the dredge-up started after reaching the asymptotic distribution (about the 20th pulse). It is difficult to calculate the AGB stars evolution and s-process nucleosynthesis. So there are few theoretical results to explain the M$\\rightarrow$S$\\rightarrow$C evolutionary sequence based on the combination of the heavier$-$lighter s-element abundances ratio and the C/O ratio, though the observational abundances exhibit a certain regularity. Zhang et al. (1998a) calculated the evolution of the surface heavy element abundances and C/O ratio for a 3$M_{\\odot}$ TP-AGB star with initial metallicity 0.015, and gave interesting results. But they adopted mean neutron exposure $\\tau_0$ and unbranch s-process path, which have been revised in these years.  In the first part of this paper, we adopt the new s-process nucleosynthesis scenario (Straniero et al. 1995; Straniero et al. 1997; Gallino et al. 1998; Busso et al. 1999 etc.), and the branch s-process nucleosynthesis path to calculate the s-process nucleosynthesis of solar metallicity 3$M_{\\odot}$ AGB stars. And then, we discuss the M$\\rightarrow$S$\\rightarrow$C sequence based on the heavy-element abundances and C/O ratio.  The importance of AGB stars nucleosynthesis is not only to explain the observational M$\\rightarrow$S$\\rightarrow$C sequence but also to be responsible for the origin of some other classes of stars with overabundances of heavy-elements.  Observations revealed that some stars with overabundances of heavy-elements were not luminous to up to the stage of AGB. Following Lambert (1991), the stars showing heavy-element overabundances are divided into two classes: (1) intrinsic TP-AGB stars$-$they include MS, S and C (N-type) stars exhibiting the unstable nucleus $~^{99}$Tc ($\\tau_{1\\over 2}=2\\times 10^5 ~$yr) as evidence that they are presently undergoing nucleosynthesis activity and the third dredge-up, and (2) extrinsic AGB stars$-$they include the various classes of G-, K-type barium stars and the cooler S, C stars where $~^{99}$Tc is not observed. It is generally believed that the extrinsic AGB stars belong to binary systems and their heavy-element overabundances come from accretion of the matter ejected by the companions (the former AGB stars, now evolved into white dwarfs) (McClure et al. 1980; Boffin \\& Jorissen 1988; Jorissen et al. 1998; Jorissen \\& Van Eck 2000; Jorissen 1999). The mass exchange took place about $1\\times 10^6$ years ago, so the $^{99}$Tc produced in the original TP-AGB stars have decayed. The accretion may either be disk accretion (Iben $\\&$ Tutukov 1985) or common envelope ejection (Paczynski 1976). Han et al. (1995) detailedly investigated the evolutionary channels for the formation of barium stars. In this paper we will only discuss the stellar wind accretion scenario because it is very important to explain the formation of barium stars (Boffin \\& Jorissen 1988; Jorissen et al. 1998).  Boffin $\\&$ Jorissen (1988) calculated qualitatively the variation of orbital elements caused by wind accretion in binary systems. They also estimated the heavy-element overabundances of barium stars. Subsequent papers (Za$\\check{c}$s 1994, Boffin $\\&$ Za$\\check{c}$s 1994) used similar methods to calculate the overabundances, and interpreted the relationship between the heavy-element abundances and the orbital periods of barium stars.  Some important conclusions have been drawn in the theory of wind accretion, but the previous calculations on orbital elements were not very reliable because of the neglect of the $\\delta r/r$ term ($r$ represents the distance between the two components of the binary system ), and using the tangential momentum conservation (Chang et al. 1997 and references therein). For the rotating binary system with wind mass loss, the total angular momentum conservation is more reasonable than tangential momentum conservation. Also, the previous calculations of heavy-element overabundances used the 'step-process' (Boffin $\\&$ Jorissen 1988, Boffin $\\&$ Za$\\check{c}$s 1994), which means that the overabundance factor changes at one single instant from 1 to $f$ ($f$ is the relative ratio of the heavy-elements to the solar abundances, and differs for different elements. Earlier calculations used the mean $f$ value of carbon stars), and then keeps the value until the end of the AGB phase. However, after the start of the TDU, the overabundance factor $f$ of the intrinsic AGB stars changes during successive dredge-ups. It is after a number of dredge-ups that the C/O ratio in the outer envelope of the intrinsic AGB star reaches the value 1, which means that the AGB star becomes a carbon star. The heavy-element overabundances of the barium star should be caused by the successive pulsed accretions and mixing.  According to the analysis to the orbital elements of barium and S stars, Jorissen \\& Mayor (1992) presented the evolutionary pathways of binaries leading to barium and S systems. They concluded that the binary systems with longer orbital period formed through wind accretion and those with shorter orbital period formed via Roche lobe overflow. But the specific range of orbital period was not presented. Jorissen et al. (1998) analyzed the orbital elements of a large sample of binary systems to give insight into the formations of barium and Tc-poor S stars. They suggested that barium stars with orbital period $P$$>$1500 days formed through wind accretion scenario. Zhang et al. (1999) suggested that the barium stars with $P$$>$1600 days formed via wind accretion according to their model. Liu et al. (2000) discuss the evolutions of orbital elements of barium stars from normal G, K giants.  Besides the orbital elements, the heavy-element abundances of barium stars have been discussed in some literatures. Busso et al. (1995) discussed the observational heavy-element abundances of barium stars. And a more detailed analysis of the abundance distributions for five stars has been performed using the method of mixing the accretion mass with the envelope mass. But the effect of mass accretion and the changes of orbital elements were not considered. Chang et al. (1997) and Liang et al. (1999) attempted to explore the relationships between the heavy-element overabundances and orbital elements of bariums stars using the binary accretion scenario, but with the shortcomings: or using the tangential momentum conservation, or adopting the old nucleosynthesis scenario of TP-AGB stars.  In the second part of this paper, we firstly reduced the variation equations of orbital elements based on the angular momentum conservation model of wind accretion, then we calculate the heavy-element overabundances of barium stars via successive pulsed accreting matter enriched heavy-elements from the intrinsic AGB stars, and mixing the matter with their envelopes.  This paper is organized as follows. The observational data of MS, S, C (N-type) stars and barium stars are given in Sect. 2. In Sect. 3, we present the model and the main parameters of AGB stars nucleosynthesis and the angular momentum conservation model of wind accretion scenario for barium stars. Sect. 4 illustrates and analyzes our results. We conclude and discuss in Sect. 5. ",
        "conclusions": "\\subsection * {5.1 The intrinsic AGB stars}  Adopting the latest AGB stars evolutionary theory and nucleosynthesis scenario, the heavy-element abundances of solar metallicity 3$M_{\\odot}$ AGB stars are calculated. It is shown that, adopting reasonable parameters, the calculated results of the heavier/lighter s-elements ratio$-$ overabundance relationships and the C/O ratio$-$overabundance relationships can fit the observations. The evolution of AGB stars along the M$\\rightarrow $S$\\rightarrow$C sequence is thus further explained from the evolutionary theories and heavy-elements nucleosynthesis of AGB stars. A 3$M_{\\odot}$ AGB star will undergo about 27 pulses before it becomes carbon star.  The results showed that the overabundances of s-process elements depend significantly on the neutron exposure: for the lower neutron exposures ($a$$<$1.5), the lighter s-process elements are more abundant than the heavier ([hs/ls]$<$0.0); for the higher neutron exposures ($a$$>$1.5), the latter can increase strongly with the increase of neutron exposures, and the lighter s-process elements change slightly ([hs/ls]$>$0.0); while the results with $a$=1.5 can explain the solar heavy-element abundance distribution, which confirm the results of Gallino et al. (1998). Moreover, $^{12}$C abundance correlates to the s-process element abundances. Both C/O ratio and s-element abundances increase with the occurrence of TDU. After some dredge-ups, C/O ratio reaches 1, from which the AGB stars become carbon stars. The observed MS, S and C (N-type) stars lie in a range of neutron exposure: $a$=0.5-2.5. The over low neutron exposure can not produce carbon stars (e.g. $a$=0.5). Our calculation thus provides a further theoretical basis for the evolution of AGB stars along the M$\\rightarrow $S$\\rightarrow$C sequence based on the heavy-element abundances and C/O ratio. The general agreements between our calculations and the observations indicate that the parameters, the theories of evolution and nucleosynthesis scenario of AGB stars adopted in calculation are reasonable.  The stars with initial main sequence mass 1.5$-$4$M_{\\odot}$ can form carbon stars (Groenewegen et al. 1995). We consider only a 3$M_{\\odot}$ star model in our calculation. So the comparison between the calculated results and the observations is somewhat rough. With further studies of the evolutionary theory of AGB stars and increasing observational data, we can expect a deeper understanding on the nucleosynthesis of AGB stars.  \\subsection * {5.2 The barium stars}  Taking the conservation of angular momentum in place of the conservation of tangential momentum for wind accretion scenario, considering the change of $\\delta r/r$ term, and not neglecting the square and higher power terms of eccentricity, the change equations of orbital elements are recalculated.  We combine wind accretion with the nucleosynthesis of intrinsic AGB stars, to calculate, in a self-consistent manner, the heavy-element overabundances of barium stars through mass accretion during successive pulsed ejection, followed by mixing.  The calculated relationships of heavy-element abundances $-$ orbital period $P$ can fit the observations within the error ranges. Moreover, the predictions of the detailed abundances of different atomic charge Z can match well the observations of 11 program barium stars with longer orbital period ($P$$>$1600 d). The corresponding neutron exposures are in the range of 0.8$<$$a$$<$2.5. The higher neutron exposure (e.g. $a$=2.0) will produce the more abundant the heavier s-elements than the lighter (see Fig. 4c); on the contrary, the abundances of the lighter s-elements are higher than the heavier with low neutron exposure (e.g. $a$=1.0, see Fig. 4a); while $a$=1.5 will produce the almost equal abundances of the heavier to the lighter s-elements (see Fig. 4h). These neutron exposures can fit the corresponding results of intrinsic AGB stars. Naturally, we can understand the observations of no-Tc MS and S stars, which are commonly believed to be the descendants of barium stars. These results of the extrinsic AGB stars confirm the reliability of the nucleosynthesis and evolution of intrinsic AGB stars. Simultaneously, the results confirm that our wind accretion model and parameters adopted are suitable.  Analyzing our results, we understand that the barium stars with longer orbital period ($P$$>$1600 d) form through wind accretion. Those with shorter orbital period ($P$$<$600 d) form through other scenarios, such as dynamically stable late case C mass transfer or common envelope ejection. Moreover, the change range of mass accretion rate should be 0.1 to 0.5 times as much as the Bondi-Hoyle's accretion rate. The corresponding range of orbital periods and mass accretion rate to the formation of barium stars still need to be tested by more observations.  At present, the orbital elements of a large sample barium stars, Tc-poor S stars have been published (Udry et al. 1998a, 1998b; Carquillat et al. 1998). But the corresponding heavy-element abundances have not been obtained. So we need the high resolution, high signal-to-noise spectral observations of these stars, which are combined with the observations of orbital elements, to research their characters and formation.  In addition, we should note that metallicity is an important factor to the AGB stars nucleosynthesis and the formation of chemical peculiar stars. Actually, $^{13}$C neutron source is related to the metallicity (Busso et al. 1995, 1999 and references therein; Gallino et al. 1999). Gallino et al. (1998) calculated the s-element nucleosynthesis of 2$M_{\\odot}$ AGB stars with low metallicity $Z$=0.01, and obtained similar abundance distribution to the 3$M_{\\odot}$ with solar metallicity AGB stars model. Busso et al. (1999) and Zhang et al. (1998b) have discussed the inverse correlation between the heavy-element abundances and the metallicity [Fe/H] of intrinsic and extrinsic AGB stars. Also, the nucleosynthesis results of low metallicity AGB stars are more suitable to study the Galactic chemical evolution in the early stage of the Galaxy. For extrinsic AGB stars, Jorissen et al. (1998) suggested that the s-process operation was more efficient in a low-metallicity population, so the Pop. II CH stars may has accreted the material much enriched in heavy elements from the former AGB companion. In this paper, our main aims are (1) calculating the AGB stars nucleosynthesis, so that we can explain the observed heavy-element abundances of MS, S and C (N-type) stars, which are near solar metallicity, and supply evidence to the M$\\rightarrow $S$\\rightarrow$C evolutionary sequence; (2) discussing the wind accretion scenario of Ba stars, which are near solar metallicity too (Za$\\check{c}$s 1994; Smith et al. 1993). So we only calculate the solar metallicity case. We will extend to study the low metallicity case in the forthcoming paper."
    },
    "0011/astro-ph0011448_arXiv.txt": {
        "abstract": "I present a series of diagrams which illustrate why the cosmic microwave background (CMB) data favor certain values for the cosmological parameters. Various methods to extract these parameters from CMB and non-CMB observations are forming an ever-tightening network of interlocking constraints. I review the increasingly precise constraints in the $\\Omega_{m} - \\Omega_{\\Lambda}$ plane and show why more cosmologists now prefer $\\Lambda$CDM cosmologies to any other leading model. ",
        "introduction": "A convenient way to interpret CMB observations is to fit the angular power spectrum of the data to parameter-dependent models. Figure \\ref{fig:data} shows the recent CMB measurements along with three such models. In Figure~\\ref{fig:binnedata}, binning of this data reduces the scatter and provides a representative region favored by the data. Important parameters that can be constrained by CMB power spectra include Hubble's constant $h$, the cosmological constant $\\Omega_{\\Lambda}$, the density of cold dark matter $\\Omega_{\\rm CDM}$, and the density of baryonic matter $\\Omega_{b}$. Figures~\\ref{fig:data}~-~\\ref{fig:zoomin} provide a qualitative feel for the lever arm that the CMB data provides for constraining these and other  parameters simultaneously. Unless stated otherwise, the models shown have the following default values: $h= 0.70$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{m}= \\Omega_{\\rm CDM} + \\Omega_{b} =0.3$, $\\Omega_{b}h^{2} = 0.020$, a power spectral index of primordial scalar density fluctuations $n_{s}=1$ and an overall normalization $Q_{10} = 18\\;\\mu$K. The grey band in Figure~\\ref{fig:binnedata} is reproduced in Figures~\\ref{fig:highh}~-~\\ref{fig:zoomin} and represents the data in a model-independent way. With it, the eye can pick out which models best fit the data.  A reduction in $h$ increases the amplitude of the first peak (Fig.~\\ref{fig:highh}). An increase in the number of baryons increases the gravitating mass of the oscillating baryon-photon fluid . This enhances the first peak (Fig.~\\ref{fig:omegabaryon}) by producing more gravitational compression as the baryons drag the photons further into the potential wells  (and further away from the potential hills). The second peak is suppressed because, before decoupling, these smaller scales experienced the same additional compression (and rarefaction) and, at decoupling, we are seeing a subdued rebound from this enhanced compression (and rarefaction), i.e., we are seeing the smaller amplitude of an oscillation whose zero level had been lowered in the previous oscillation by the effect of additional baryons. An increase in $\\Omega_{m}$ decreases the amplitude of the first peak (Fig.~\\ref{fig:breakdegeneracy}). For discussion of the physics of the parameter dependencies of features in the CMB power spectrum see e.g. Hu \\& Sugiyama (1995), Hu (1995), Tegmark (1996), Lineweaver et al. (1997).  \\clearpage \\begin{figure}[!ht] \\plotfiddle{cl1.ps}{7.0cm}{0}{63}{53}{-187}{-57} \\caption{{\\footnotesize Compilation of recent CMB observations from 22 experiments compared to three popular cosmological models. At the last IAU, all three models were serious contenders for reality. This is no longer the case.}} \\label{fig:data} \\end{figure} \\begin{figure}[!hb] \\plotfiddle{cl1fill.ps}{7.0cm}{0}{63}{53}{-187}{-57} \\caption{{\\footnotesize Same as Fig.~\\ref{fig:data}, but the data points have been binned and averaged and rebinned multiple times to provide a series of points and a grey area defined by the 1 $\\sigma$ error bars of the binned points. The points are not independent. The well defined position of the first peak in the spectrum at $\\ell_{p} \\sim 220$ easily rules out the $(\\Omega_{m}, \\Omega_{\\Lambda}) = (0.3,0.0)$ model shown which peaks at $\\ell_{p} \\sim 400$ (see e.g. Lineweaver 1998).}} \\label{fig:binnedata} \\end{figure}  \\clearpage \\begin{figure}[!ht] \\plotfiddle{cl1h.ps}{6.3cm}{0}{60}{48}{-175}{-52} \\caption{{\\footnotesize {\\bf Standard CDM is consistent with CMB data, but only if $h$ is extremely low.} Until recently, the standard cold dark matter model ($\\Omega_{m}, \\Omega_{\\Lambda}) = (1,0)$ was the leading cosmological candidate, but the amplitude of the first peak is too low in these models unless $h \\sim 0.40$. See the $h$ contours in Panel A of Figure~\\ref{fig:science} for more details. The grey region is preferred by the data and comes from the previous figure.}} \\label{fig:highh} \\end{figure} \\begin{figure}[!hb] \\plotfiddle{cl1obh2.ps}{6.3cm}{0}{60}{48}{-175}{-52}   % \\caption{{\\footnotesize {\\bf Why does the CMB prefer high values of $\\Omega_{b} h^{2}$?} Compared to big-bang-nucleosynthesis-dependent estimates ($\\Omega_{b} h^{2} \\approx 0.020$ e.g. Tytler et al. 2001) the first peak of the CMB data has a slightly higher amplitude and the second peak has a slightly lower amplitude. Raising $\\Omega_{b} h^{2}$ to 0.030 fits the data better because it raises the first peak and lowers the second peak. More precise data in the $\\ell \\sim 500$ region will soon be available to help solve this tentative conflict. }} \\label{fig:omegabaryon} \\end{figure}  \\clearpage \\begin{figure}[!ht] \\plotfiddle{cl1flath455565.ps}{6.5cm}{0}{60}{48}{-187}{-57}       % \\caption{{\\footnotesize {\\bf Degeneracy is a big problem.} These three flat models ($\\Omega_{tot} = \\Omega_{m} + \\Omega_{\\Lambda} = 1, \\Omega_{k} = 0$) have very different values for $\\Omega_{m}$ and $\\Omega_{\\Lambda}$, but the CMB cannot differentiate between them if $h$ can vary . Reducing $h$ raises the peak (Fig.~\\ref{fig:highh}) while increasing $\\Omega_{m}$ lowers the peak. Similar but usually more subtle degeneracies become more numerous as the dimensionality of explored parameter space increases. An important issue is what range of values does one allow for $h$ and other parameters.}} \\label{fig:degeneracy} \\end{figure} \\begin{figure}[!hb] \\plotfiddle{cl1flath70.ps}{6.5cm}{0}{60}{48}{-187}{-57}      % \\caption{{\\footnotesize {\\bf Degeneracies can be broken.} Precise non-CMB constraints (such as an independent measurement of $h$) or much more precise CMB data can break degeneracies. Here, setting $h = 0.70$ breaks the degeneracy of Fig. 5, allowing the CMB to discriminate between various flat models.}} \\label{fig:breakdegeneracy} \\end{figure}  \\clearpage \\begin{figure}[!ht] \\plotfiddle{cl1omegakv2.ps}{5.8cm}{0}{57}{45}{-175}{-52} \\caption{{\\footnotesize {\\bf The CMB can measure the geometry of the Universe, $\\Omega_{k}$, better than other cosmological observations.} The position of the first peak in the CMB data, $\\ell_{p}$, is probably the best measurement we have of the geometry of the universe. It is an excellent tool to answer the question: `Is the universe spatially open, flat or closed?' Or, in terms of cosmological parameters ($\\Omega_{k} = 1- \\Omega_{tot}$), `Is $\\Omega_{k}$ greater than, equal to or less than zero?' (or equivalently, `Is $\\Omega_{tot}$ less than, equal to or greater than one?') We have set $\\Omega_{k} = 0.0, 0.2, 0.4$ (flat, open, more open, respectively). The vertical lines indicate the peak in the data (Fig.~\\ref{fig:zoomin}).}} \\label{fig:geometry} \\end{figure} \\begin{figure}[!hb] \\plotfiddle{cl1lpeakv4.ps}{5.8cm}{0}{57}{45}{-175}{-52} \\caption{{\\footnotesize {\\bf Where is the peak?} This is a blow-up of the region of the first peak with the multiply-binned points from Fig.~\\ref{fig:binnedata} and the models from Fig.~\\ref{fig:omegabaryon}. An eye-ball model-independent estimate of the position of the first peak is $\\ell_{p} \\approx 225 \\pm 25$, indicated by the vertical lines (same as in previous figure). This result is based on the data shown in Fig.~\\ref{fig:data}. The points are {\\it not} independent.}} \\label{fig:zoomin} \\end{figure}  \\clearpage  Just as the SNIa data is our strongest lever arm for determining the acceleration of the universe, $q_{o}$, the CMB data is our strongest lever arm for determining the geometry of the universe, $\\Omega_{k}$. The ability of the CMB to constrain $\\Omega_{k}$ can be seen in Fig.~\\ref{fig:science}, panel A and Fig.~\\ref{fig:latest} in which the CMB contours are elongated in the $\\Omega_{k} = $ {\\it constant} direction (upper left to lower right). However, since both $\\Omega_{k}$ and $h$ control the position of the peak, there is a slight degeneracy. It is difficult to separate the effect of the spatial geometry of the universe from the effect of $h$. This degeneracy is reflected in the width of the elongated CMB constraints in the $\\Omega_{m} - \\Omega_{\\Lambda}$ plane. The models in Fig.~\\ref{fig:zoomin} have $\\Omega_{k} = 0$, $h = 0.70$. The data and the high baryon model peak at $\\ell \\approx 225$. This crude eye-ball estimate should be compared to the more careful but model-dependent estimates of Bond et al. (2001, $\\ell_{p} = 212 \\pm 7$) and Hu et al. (2001, $\\ell_{p} < 218$ at the $2\\sigma$ level based on the Boomerang and Maxima results only). ",
        "conclusions": ""
    },
    "0011/astro-ph0011162_arXiv.txt": {
        "abstract": "Hipparcos proper motions and trigonometric parallaxes allow the derivation of secular parallaxes which fix the distances to individual stars in the Hyades cluster to an accuracy of $\\sim$2~percent. The resulting color-absolute magnitude diagram for 92 high-fidelity single members of the cluster displays a very narrow main sequence, with two turn-offs and associated gaps. These occur at the locations where the onset of surface convection affects the $B-V$ colors, as predicted by B\\\"ohm--Vitense thirty years ago. The new distances provide stringent constraints on the transformations of colors and absolute magnitudes to effective temperatures and luminosities, and on models of stellar interiors. ",
        "introduction": "The Hyades open cluster is one of the key calibrators of the absolute magnitude-spectral type relation and the mass-luminosity relation. The small distance to the cluster ($\\sim$45~pc) has a number of advantages for studies of its Hertzsprung--Russell diagram: (i) the foreground interstellar reddening and extinction is negligible ($E(B-V) = 0.003 \\pm 0.002$~mag \\citep[e.g.,][]{tay1980}), (ii) the large proper motion ($\\mu \\sim 111$~mas~yr$^{-1}$) and peculiar space motion ($\\sim$35~km~s$^{-1}$) facilitate proper motion- and radial velocity-based membership determinations, and (iii) the stellar content can be probed to low masses relatively easily.  Among the $\\sim$400 known members are white dwarfs, red giants, mid-A stars in the turn-off region, and main-sequence stars down to at least $\\sim$$0.10~M_\\odot$ M dwarfs \\citep*{bhj1994}. At the same time, the proximity of the Hyades complicates astrophysical interpretation: the tidal radius of $\\sim$10~pc corresponds to a significant depth along the line of sight. As a result, the precise definition and location of the main sequence and turn-off region in the Hertzsprung--Russell diagram, and the accuracy of the determination of, e.g., the helium content and age of the cluster, have always been limited by the measurement error of the distances to individual stars. Even Hipparcos parallax uncertainties translate into absolute magnitude errors of $\\ga$0.10~mag at the mean distance of the cluster, whereas $V$-band photometric errors only account for $\\la$0.01~mag uncertainty for most members.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=0.0, width=18.0truecm, clip=true, keepaspectratio=true] {debruijne01.eps} \\end{center} \\caption[]{Color-absolute magnitude diagram for 92 high-fidelity members of the Hyades cluster. This sample excludes all members beyond 40~pc from the cluster center, double and multiple stars, and stars with suspect secular parallaxes (see BHZ for details). The absolute magnitudes were computed using the observed $V$-band magnitudes and secular parallaxes derived by \\citet{bhz2000}. The $B-V$ colors were taken from the Hipparcos Catalog. The arrows indicate the two B\\\"ohm--Vitense turn-offs and associated gaps, which are most likely caused by sudden changes in the properties of convective atmospheres. The gaps between $B-V=0.15$ and $0.20$, between $B-V=0.30$ and $0.35$, and the gap around $B-V=0.95$ mag are caused by the suppression of double, multiple, and peculiar stars from our sample (cf.\\ figure 21 in \\citet{per1998}; see \\S 2); the region between $B-V=0.30$ and $B-V=0.35$ mag, e.g., is occupied by Am-type stars, which have a high incidence of duplicity \\citep[e.g.,][]{jj1987}.} \\label{fig:camd} \\end{figure*}  It was realized long ago \\citep[e.g.,][]{bos1908} that the small internal velocity dispersion of the Hyades makes it possible to improve the trigonometric parallaxes of individual member stars by kinematic modelling of their proper motions. Various recent studies derived such secular parallaxes for Hyades members from Hipparcos astrometry \\citep*{per1998,bru1999,dlm1999,ng1999}, but none investigated the resulting Hertzsprung--Russell diagram in any detail. Yet, as the relative proper motion accuracy is effectively three times larger than the relative trigonometric parallax accuracy, the secular parallaxes are three times more precise than the trigonometric parallaxes, and provide a unique opportunity to obtain a well-defined and absolutely calibrated Hertzsprung--Russell diagram for an open cluster. For this reason we have redetermined the secular parallaxes for the Hyades, using the comprehensive procedure described by \\citet{bru1999} \\citep*[cf., e.g.,][]{lmd2000}. We describe the full derivation of the secular parallaxes elsewhere \\citep*[BHZ]{bhz2000}. This includes a new analysis of the space motion and internal velocity dispersion of the Hyades, a validation of the secular parallaxes, with a detailed investigation of the effect of velocity structure in the cluster and of the presence of small-angular-scale correlations in the Hipparcos data, and an extensive analysis of the construction of the physical Hertzsprung--Russell diagram $(\\log T_{\\rm eff}, \\log L)$.  In the course of this work, we constructed a color-absolute magnitude diagram for a sample of high-fidelity single members of the Hyades, and discovered that the improved accuracy reveals two turn-offs and associated gaps in the main sequence at the locations predicted by B\\\"ohm--Vitense, in the seventies. These are the topic of this Letter.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=-90.0, width=8.0truecm, clip=true, keepaspectratio=true] {debruijne02.eps} \\end{center} \\caption[]{Schematic color-absolute magnitude diagram of an open cluster with a B\\\"ohm-Vitense gap. Stars left of the dashed vertical line, at $B-V\\sim 0.3$ mag, have radiative envelopes. Surface convection sets in at the line, and all stars to the right of it have convective envelopes. While the onset of convection does not alter luminosity or effective temperature, it causes redder colours because of lower temperatures in deeper layers (see \\S 3). The resulting color-absolute magnitude diagram therefore shows a turn-off and gap.} \\label{fig:bvgap} \\end{figure*} ",
        "conclusions": "Narrow main sequences in color-magnitude diagrams are readily observable for distant clusters, but the absolute calibration of the Hertzsprung--Russell diagram of such groups is often uncertain due to their poorly determined distances and the effects of interstellar reddening and extinction. The latter problems are alleviated significantly for nearby clusters, but at the price of introducing a considerable spread in the location of individual members in the Hertzsprung--Russell diagram as a result of their resolved intrinsic depths, and relatively poorly determined individual distances. We have shown here that the Hyades are unique in that the secular parallaxes derived from the Hipparcos astrometry for the members are sufficiently accurate to calibrate the main sequence of this nearest open cluster in absolute terms.  The color-absolute magnitude diagram for the Hyades shown in Figure~\\ref{fig:camd} reveals two turn-offs and associated gaps in the main sequence. We identify these with the so-called B\\\"ohm--Vitense gaps, which are most likely related to sudden changes in the properties of surface convection zones in the atmospheres of stars with $B-V \\sim 0.30$ and $\\sim$0.45~mag, which significantly affect the emergent UV and blue-optical fluxes, and thus the $B-V$ color. We show in BHZ that this substructure in the $(B-V, M_V)$ diagram provides a strong constraint on stellar models, requiring an improvement in the available transformations from colors and absolute magnitudes to effective temperatures and luminosities.  The future astrometric satellites FAME and GAIA will improve the accuracy of stellar astrometry into the micro-arcsecond regime. They will make it possible to provide accurate membership and high-precision absolutely-calibrated main sequences for all star clusters and associations to distances of at least 2 kpc. This will provide the ability to test stellar models over a range of metallicities and ages, and may well reveal further substructure."
    },
    "0011/astro-ph0011481_arXiv.txt": {
        "abstract": "We present a method, based on the kinematic analysis of the Galactic disk stars, to clarify whether the internal motions of the stellar system in spiral arms follow those expected in the density wave theory. The method relies on the comparison with the linear relation between the phases of spatial positions and epicyclic motions of stars, as drawn from the theory. The application of the method to the 78 Galactic Cepheids near the Sun, for which accurate proper motions are available from the {\\it Hipparcos} Catalogue, has revealed that these Cepheids hold no correlation between both phases, thereby implying that their motions are in contradiction with the theoretical predictions. Possible reasons for this discrepancy are discussed and future prospects are outlined. ",
        "introduction": "Spiral structures of galaxies have been studied for a long time in order to understand how these structures are formed (e.g., Roberts, Roberts \\& Shu 1975; Rohlfs 1977; Binney \\& Tremaine 1987). One of the proposed models to explain spiral arms is that they are just material arms, where the stars originally making up a spiral arm remains in the arm even at the later time. However this simple model holds a wellknown problem which is called ``winding problem'': the differential rotation in galactic disks winds up the arm in a short time compared with the age of galaxies, so that the spiral pattern would be too tightly wound compared with the observed spiral structures. In contrast, the currently most popular model, which is free from the winding problem, is the density wave theory (Lin \\& Shu 1964), where a spiral arm is regarded as a wave and wavelike oscillation of stellar motions propagates through galactic disks. In this picture, the global spiral pattern is sustained independently of individual stars moving at different angular velocities. For a comprehensive review of the density wave theory, see, e.g., Rohlfs (1977) and Binney \\& Tremaine (1987).  The density wave theory has been suggested by various observational aspects in spiral galaxies, including the relative distributions of dust lane, interstellar gas, and H II regions across the arms (Fujimoto 1968; Roberts 1969; Rohlfs 1977), intensity distribution of radio continuum radiation (Mathewson, van der Kruit, \\& Brouw 1972), and systematic variation of gaseous velocity fields near the arms (e.g., Visser 1980). In particular, recent high-resolution observations using CO emission have revealed detailed streaming motions of molecular gas, which are generally in agreement with predictions of the density wave models (e.g., Kuno \\& Nakai 1997; Aalto et al. 1999). However, we note that these observational results provide only an outcome of nonlinear interaction between interstellar matter and background stellar arms, and it is yet unknown whether the motions of stars {\\it themselves}, which make up spiral pattern, actually follow those predicted by the density wave theory. In this regard, the direct access to detailed stellar kinematics in disks is possible only in our Galaxy.  Here, we present a method to clarify this issue, based on the analysis of local kinematics of disk stars. We then apply the method to 78 Cepheids in the solar neighborhood, for which the precise data of proper motions are available from the {\\it Hipparcos} Catalogue (ESA 1997). Also, the distances to these sample stars can be accurately estimated from the period-luminosity relation, so that combined with the radial velocity data, the full three-dimensional velocities are available. We note here that although we focus on the local kinematics of spiral arms in this work, the method we develop here can be applied to the motions of more remote stars distributed over a whole disk, for which precise astrometric data will be provided by the next-generation satellites such as {\\it FAME} and {\\it GAIA}.  Our paper is organized as follows. In \\S 2, we describe the method to determine whether or not the motions of stars agree with those expected in the density wave theory. In \\S 3, we show the detail of the sample stars and the fundamental parameters of our Galaxy adopted in this work. The application of our method to the sample stars is shown in \\S 4. Finally, \\S 5 is devoted to discussion and conclusions. ",
        "conclusions": "We have presented a method, based on the analysis of local kinematics of disk stars, to clarify whether the motions of the stars, which make up the spiral arms, follow those expected in the density wave theory. The method utilizes the comparison with the expected linear relation between the ``position phases'' of the stars $\\chi$ and those of their epicyclic motions $\\phi$, as given in equation (\\ref{2}). The application of the method to the 78 Galactic Cepheids within 4 kpc from the Sun, for which accurate proper motions are available from {\\it Hipparcos}, has revealed that the relation between $\\chi$ and $\\phi$ for the Cepheids does not show the expected linear relation. Based on the quantitative analysis using the deviation $\\Delta^2$, we conclude that the observed motions of the Cepheids are well reproduced by the random model having no correlation between $\\chi$ and $\\phi$.  There are a couple of possibilities to explain the current results, even if the spiral arms follow the density wave theory. First, the spiral structure around the Sun may not be simple as given in equation (\\ref{1}). In fact, many of our sample Cepheids belong to the local spiral structure called the Orion arm, for which the definite conclusion on its spatial structure is yet to be reached. It is frequently expressed as ``Orion spur'', having a rather irregular pattern compared to other large-scale arms, Sagittarius and Perseus arm (see, e.g., Gilmore, King, \\& van der Kruit 1989). The existence of the Orion arm may give disturbances on the density wave motions of stars induced by these large-scale arms. Second, the Cepheids we have adopted here may still convey systematic velocities of dens gas clouds from which these stars were formed, in the form of the streaming motions. If there still exist some individual streaming motions among the sample stars, such motions may violate the ideal linear relation between $\\chi$ and $\\phi$ expected for the density wave motions. Third, the Cepheids we have adopted here may have already experienced some scattering by dens gas clouds, thus having large velocity dispersions (Spitzer \\& Schwarzschild 1953). However, our experiment in \\S 4 (dotted line in Figure 4) implies that the effect of velocity dispersions of our sample on the result appears to be minor.  In order to settle the last issue described above more clearly, we have repeated our analysis using younger populations with smaller velocity dispersions than the Cepheids. As such young stars, we have adopted the O-B5 stars, although due to their fainter luminosities than Cepheids, the sample with available proper motions is confined to the narrower region near the Sun. These sample stars are taken from the NASA SKY2000 Master Star Catalog Ver.~2 (Sande et al. 1998) which provides almost 300,000 stars brighter than 8 mag. The catalog contains many basic quantities, such as MK classification, luminosity class, apparent magnitude, color, radial velocity, and so on. We have then calibrated distances using {\\it Hipparcos} parallaxes or spectroscopic distances using the program kindly supplied by Drs. M.~S\\^oma and M.~Yoshizawa, and also obtained accurate proper motions by the cross-identification with the {\\it Hipparcos} and ACT Reference Catalogs (Urban, Corbin, \\& Wycoff 1998). After removing binaries and multiples, we have selected 773 O-B5 stars for which full three-dimensional velocities are available. Then, the application of the method we have developed here has revealed that the deviation $\\Delta^2$ as a function of $m$ remains essentially constant of the order of 3000, without showing any noticeable minimum at a specific value of $m$. Thus, even the motions of such young populations with small velocity dispersions are in contradiction with the density wave theory. We note here that most of these O-B5 stars are located within $\\sim 1$ kpc from the Sun, so the effect of the local irregular spiral on the result cannot be negligible.  More definite conclusions on the issue we have addressed here require the assembly and analysis of much larger numbers of stars with accurate distances and proper motions, so that the statistical fluctuation in the result can be significantly reduced. Also, it is necessary to assemble the data of more remote stars over a large fraction of the disk, thereby diminishing effects of local irregular spiral structures on the kinematic analysis. Indeed, next-generation satellites such as {\\it FAME} and {\\it GAIA} will provide very precise astrometric data for huge numbers of the Galactic stars, and will thus offer us an opportunity to assess detailed motions of disk stars in conjunction with the density wave theory."
    },
    "0011/astro-ph0011354_arXiv.txt": {
        "abstract": "Oscillation frequencies are the most accurate properties one can measure for a star, potentially allowing detailed tests of stellar models and evolution theories.  We briefly review asteroseismology for two classes of stars.  In $\\delta$~Scuti variables, the main problems are the identification of the observed modes and the theoretical treatment of rotation.  In solar-like stars the main difficulty is the tiny amplitudes, but credible detections are now being made.  These confirm that stars are oscillating at the approximately the expected levels, but suggest that amplitudes scale as $1/g$ rather than $L/M$.  We also stress the importance of multi-site campaigns.  Several space missions will be launched over the coming years, promising an exciting future for asteroseismology. ",
        "introduction": "Asteroseismology involves an interplay between observations of stellar oscillations and theoretical model calculations.  It can be done when a set of oscillation frequencies is known for a given star and, at the same time, a set of theoretical model frequencies can be calculated.  The motivation for doing asteroseismology is that oscillation frequencies are the most accurate properties one can measure for a star and we may therefore, at least in principle, be able to perform a detailed test of stellar modelling and evolution theories.  This potentially very promising tool has motivated a huge observational effort with the aim of obtaining accurate oscillation frequencies.  At the same time, a substantial amount of work has been put into improvements of stellar modelling with the goal of being able to fit model frequencies to the observed oscillation frequencies.  In helioseismology, as can be seen from the papers presented in these proceedings, most current work concentrates on the details.  We are already able to measure accurately many important properties of the solar interior. Unlike asteroseismology, helioseismology is founded on an immense amount of high-quality data supported by an equally detailed set of state-of-the-art models.  It is impossible to imagine that asteroseismology will ever reach a level similar to that in which we find helioseismology today, concerning analysis techniques, data quality and the level of the results.  The reasons for this lie in the main differences between the two subjects: \\begin{itemize}\\itemsep=0pt  \\vspace*{-2ex}  \\item Asteroseismology works on distant stars whose basic properties will always be less well determined than those of the Sun.  These include age, composition, radius, mass, atmospheric properties and neutrino flux.  The uncertainties will affect the quality of the calculated stellar models. One example is age, which is known for the Sun from radioactive dating of the solar system.  To illustrate the importance of this parameter for the current solar model, consider the following simple question: What age would we determine for a solar model -- based on the observed properties of the Sun -- if we did not already know the age of the solar system?  \\item The surfaces of stars are essentially unresolved, which limits asteroseismology to modes of low degree.  Much of what we have learned about the Sun is based on high degree modes.  \\item The Sun is a relatively simple system.  This makes it very interesting, since we may have a chance to understand it!  Many stars seem to be much more complicated.  Hopefully we will find some important physical properties that are not known at present, which may turn out to be important for understanding the details of those stars.  \\item There is only one Sun, while there are billions of other stars.  Many people work on the details in helioseismology, but we can't expect several billion astronomers to be working in asteroseismology!  \\item In helioseismology we see some big networks and dedicated telescopes (GONG, IRIS, BiSON), and we have many years of space research (SOHO, SolarMAX, IPHIR).  In asteroseismology we find many smaller campaigns and networks, but very few dedicated telescopes.  The space projects are just beginning.  So far, there is a tremendous lack of high-quality data.  \\item Asteroseismology works on fainter stars and so has lower sensitivity than helioseismic observations.  One will therefore generally be restricted to oscillations with relatively high amplitudes.  \\end{itemize}  To explain why asteroseismology is so far behind helioseismology, we could identify three important elements from the above list: \\begin{itemize}\\itemsep=0pt  \\vspace*{-2ex} \\item The imprecision of basic stellar properties. \\item The complication of the stellar physics. \\item The lack of high-quality oscillation data. \\end{itemize} ",
        "conclusions": "In this review we have described the current status of asteroseismology for A, F and G subgiants and main sequence stars.  The main conclusions are as follows: \\begin{itemize}  \\item Asteroseismology is far from helioseismology in terms of techniques and results.  \\item There are plenty of challenges for theoreticians, such as mixing, fluid motions, turbulent convection and deviations from spherical symmetry. The most important area to focus on is rotation.  \\item Asteroseismology is observationally driven.  Ground-based velocity observations are now achieving believable detections, and it is time for multi-site campaigns.  A large ground-based network is very desirable.  \\item A number of exciting space missions are in various stages of planning and construction (MOST, COROT, MONS and Eddington).  If these succeed, it seems likely that asteroseismology will enter a golden age.  \\end{itemize}"
    },
    "0011/astro-ph0011024_arXiv.txt": {
        "abstract": "We present new optical and ultraviolet spectroscopy of the anomalous EUV emitter HD~199143 (F8V). High resolution spectra in the H$\\alpha$ and Na\\,{\\sc i}\\,D wavelength regions show evidence for very rapid (a few hundred km~s$^{-1}$) rotation of the stellar photosphere. Using archive {\\it IRAS} data we also show that the star has excess emission above photospheric levels at 12~$\\mu$m. {\\it IUE} data of HD~199143 reveal the presence of emission lines of Mg\\,{\\sc ii}, C\\,{\\sc i}, C\\,{\\sc ii}, C\\,{\\sc iii}, C\\,{\\sc iv}, Si\\,{\\sc iv}, He\\,{\\sc ii} and N\\,{\\sc v} and show a large variability, both in the continuum and line fluxes. We propose that all available data of HD~199143 can be explained by assuming that is has been spun up by accretion of material from a close T~Tauri like companion, responsible for the emission lines, the ultraviolet variability and the excess infrared emission. The bursting or flaring nature of this object, mostly in high energies, could be explained as episodic mass transfer between the star and its close companion. We show that HD~199143 and the Li-rich late-type dwarf BD$-$17\\degr6128 form a physical pair and suggest that both may be part of a new nearby (48~pc) young ($\\sim$ 10$^7$ yr) stellar association in Capricornius. ",
        "introduction": "Zuckerman \\& Webb (2000) sketch a picture of the recent star formation history of the solar neighbourhood in which 10--40 million years ago an ensemble of molecular clouds were forming stars at a modest rate near the present position of the Sun. About 10~Myrs ago, the most massive of these newly formed stars exploded as a supernova, terminating the star formation episode and generating the very low density region seen in most directions from the present Sun (Welsh et al. 1998). This scenario can not only explain the presence of young stellar groups close to the earth, but also explains how the $\\beta$~Pic moving group can be so young (20~Myr; Barrado y Navascu\\'es et al. 1999), and yet so close. However, currently this scenario is largely speculative.  HD~199143 is a poorly studied bright ($V$ = 7\\fm27) star in the constellation of Capricornius. It has been classified as F8V in the Michigan Spectral Survey (Houk \\& Smith-Moore 1988), after an initial classification of G0 by Cannon \\& Mayall (1949). The star would be completely inconspicuous, if it hadn't been detected as a bright extreme-ultraviolet source by the {\\it ROSAT} (2RE J205547$-$170622) and {\\it Extreme Ultraviolet Explorer} (2EUVE J2055$-$17.1) missions (Pounds et al. 1993; Malina et al. 1994; Pye et al. 1995; Bowyer et al. 1996).  In this {\\it Letter} we present new optical and ultraviolet spectroscopy of HD~199143 and show that it is a variable and rapidly rotating F8V star. We infer that all characteristics of the HD~199143 system can be explained by assuming that it is a binary system, in which the primary has been spun up by accretion of mass from a low-mass companion. Its association with a previously studied T~Tauri-like system (BD$-$17\\degr6128) suggests that these two stars could be the first two members of a close (48~pc) new region of recent star formation and may provide compelling support for the star formation history of the solar neighbourhood outlined in the first paragraph. ",
        "conclusions": "The presence of a normal late-type companion could not explain the ultraviolet excess, or the {\\it EUVE} and {\\it ROSAT} detections of HD~199143. However, the presence of an accretion disk around our hypothetical companion, such as that found in LMXB or T~Tauri systems, might easily explain those properties, as well as the infrared excess, the presence of emission lines and the variability. In such a scenario, the high rotational velocity of HD~199143 could be due to a spin-up in its past by accretion from the companion.  At first glance, a scenario in which a nearby main-sequence star like HD~199143 would have a T~Tauri-like companion would seem far-fetched. However, Mathioudakis et al. (1995) report the presence a strongly flaring K7e--M0e dwarf with a high Li abundance only 5 arcminutes from HD~199143. The optical spectrum of this star, BD$-$17\\degr6128, is identical to that of many T~Tauri stars. From Digital Sky Survey images we identify BD$-$17\\degr6128 with HD~358623. An inspection of the Tycho-2 Catalogue (H{\\o}g et al. 2000) shows that this star has a proper motion of 59 $\\pm$ 3 and $-$63 $\\pm$ 3 mas~yr$^{-1}$ in $\\mu_\\alpha$ and $\\mu_\\delta$, identical to that of HD~199143. From the fact that HD~358623 is the only star within a 5 degree radius for which this is the case, we exclude the possibility that this could be a coincidence and conclude that the two stars form a genuine proper motion pair. Using the data by Mathioudakis et al. (1995), and the newly determined distance, we compute the absolute luminosity of BD$-$17\\degr6128 to be 0.34 $\\pm$ 0.06 L$_\\odot$, employing a similar procedure to that followed for HD~199143. Comparison with the pre-main sequence evolutionary tracks by D'Antona \\& Mazzitelli (1997) yields an age of 10$^7$ years for BD$-$17\\degr6128, consistent with a T~Tauri nature of this star.  Using the radial velocity of HD~199143 determined in Section~2, and the parallax and proper motions listed in the {\\it Hipparcos} catalogue, we compute the galactic space velocity components $(U, V, W)$ of HD~199143 to be $(-10 \\pm 13, -13 \\pm 6, -13 \\pm 6)$ km~s$^{-1}$.  This space motion is similar to that of many stars in the vicinity of the Tucanae and TW Hydra associations (Zuckerman \\& Webb 2000), suggesting that these stars might have formed from the same cloud complex.  We conclude that HD~199143 and BD$-$17\\degr6128 could very well be the first two members of a region of recent star formation similar to the TW Hydrae Association and the newly identified Tucanae Association (Kastner et al. 1997; Zuckerman \\& Webb 2000). If confirmed, a further study of these two enigmatic stars could lead to a better understanding of the star formation history in the solar neighbourhood."
    },
    "0011/astro-ph0011212_arXiv.txt": {
        "abstract": "We create mock pencil-beam redshift surveys from very large cosmological $N$-body simulations of two Cold Dark Matter cosmogonies, an Einstein-de Sitter model ($\\tau$CDM) and a flat model with $\\Omega_0 =0.3$ and a cosmological constant ($\\Lambda$CDM). We use these to assess the significance of the apparent periodicity discovered by Broadhurst \\etal (1990). Simulation particles are tagged as `galaxies' so as to reproduce observed present-day correlations. They are then identified along the past light-cones of hypothetical observers to create mock catalogues with the geometry and the distance distribution of the Broadhurst \\etal  data. We produce 1936 (2625) quasi-independent catalogues from our $\\tau$CDM ($\\Lambda$CDM) simulation. A couple of large clumps in a catalogue can produce a high peak at low wavenumbers in the corresponding one-dimensional power spectrum, without any apparent large-scale periodicity in the original redshift histogram. Although the simulated redshift histograms frequently display regularly spaced clumps, the spacing of these clumps varies between catalogues and there is no `preferred' period over our many realisations. We find only a 0.72 (0.49) per cent chance that the highest peak in the power spectrum of a $\\tau$CDM ($\\Lambda$CDM) catalogue has a peak-to-noise ratio higher than that in the Broadhurst \\etal data. None of the simulated catalogues with such high peaks shows coherently spaced clumps with a significance as high as that of the real data. We conclude that in CDM universes, the regularity on a scale of $\\sim 130h^{-1}$Mpc observed by Broadhurst \\etal has {\\it a priori} probability well below $10^{-3}$. ",
        "introduction": "The redshift distribution of galaxies in the pencil-beam survey of Broadhurst et al. (1990, hereafter BEKS) displayed a striking periodicity on a scale of 128$h^{-1}$Mpc. This result has attracted a good deal of interest over the subsequent decade, and the significance and nature of periodicity or regularity in the distribution of galaxies has remained the subject of a stimulating debate in both observational and theoretical cosmology. Although a number of studies have been devoted to the BEKS pencil-beam survey and other similar surveys, several fundamental questions remain unanswered.  From the theoretical viewpoint, it is important to decide whether such apparently periodic galaxy distributions can occur with reasonable probability in a Cold Dark Matter universe, or require physics beyond the standard paradigm. Performing large simulations can directly address this question. The first simulation specifically designed for pencil-beam comparisons was that of Park and Gott (1991, hereafter PG). Their rod-shaped CDM simulation allowed them to create twelve quasi-independent mock pencil-beam surveys similar in length to that of BEKS. One of their samples appeared `more periodic' than the BEKS data according to the particular statistical test they used for comparison. Other authors (Kurki-Suonio et al.; Pierre 1990; Coles 1990; van de Weygaert 1991; SubbaRao \\&  Szalay 1992) have used purely geometrical models such as cubic lattices and Voronoi foams to explore the implications of apparent regularities similar to those found by BEKS. In particular, SubbaRao and Szalay (1992) presented a sequence of Monte Carlo simulations of surveys of Voronoi foams, showing that such a model can successfully reproduce the data as judged by a variety of statistical measures, for example, the heights, positions and signal-to-noise ratios of the highest peaks in the power spectra. Kaiser \\& Peacock (1991) argued that the highest such peaks in the BEKS data are not sufficiently significant to be unexpected in a CDM universe, but did not support this conclusion with detailed simulations. Dekel et al. (1992) introduced other statistics, more similar to those of PG, and again concluded that the apparent periodicity seen in the real data is not particularly unlikely in any of the toy models they used for comparison. Their models include Gaussian models with an extreme initial power spectrum with power only on scales $\\sim 100h^{-1}$Mpc. They found regular `galaxy' distributions a few per cent of the time and concluded that the BEKS data do {\\it not} rule out all Gaussian models. However, these theoretical studies did not give any clear answer to the question posed above: is the BEKS regularity compatible with the standard CDM paradigm? We attempt to answer this below using versions of all the statistical tools developed in earlier papers.  There have been several interesting observational developments after BEKS. Willmer et al.  (1994) found that, if the original BEKS deep survey at the North Galactic Pole had been carried out 1 degree or more to the west, many of the peaks would have been missed. On the other hand, Koo et al.  (1993) added new data from a wider survey to the original BEKS data and found the highest peak in the power spectrum to be further enhanced. They also analysed another set of deep pencil-beam surveys and found a peak of weaker significance on the {\\it same} scale, 128 $h^{-1}$Mpc. This raises another question: is 128 $h^{-1}$Mpc a preferred length scale for the galaxy distribution?  Further support for such a preferred scale has been presented by Tully et al.  (1992), Ettori et al.  (1997) and Einasto et al. (1997). Thus one can wonder whether a single scale could be indicated with such apparent consistency within the CDM paradigm.  With the important exception of the work of PG there has been surprisingly little comparison of the BEKS data with direct simulations of standard CDM cosmogonies. Even before the BEKS discovery, White et al.  (1987) had shown that pencil-beams drilled through periodic replications of their CDM simulations frequently showed a kind of `picket fence' regularity in their redshift distribution. Frenk (1991) confirmed this result and concluded that regular patterns similar to that seen in the BEKS data are easy to find in their simulations. However, it is clearly dangerous to make use of periodic replications of a simulation when assessing the significance of apparent periodicities in the redshift distribution. It is preferable to simulate a volume large enough to encompass the whole survey. Furthermore, since many independent artificial surveys are needed to establish that the real data are highly unlikely in the cosmogony simulated, the simulated volume must be fully three-dimensional (unlike that of PG) to allow the creation of many quasi-independent lines-of-sight. A final consideration is that the BEKS data reach to redshifts beyond 0.3, so that evolution of clustering along the survey may not be negligible.  In this paper we investigate the distribution of `galaxies' along the past light-cones of hypothetical observers. Particle positions and velocities on these light-cones were generated as output from the Hubble Volume Simulations (Evrard et al.  2000). These very large CDM $N$-body simulations were recently performed by the Virgo consortium and each used $10^9$ particles to follow the evolution of the matter distribution within cubic regions of an $\\Omega=1$ $\\tau$CDM ($\\Omega=0.3$ $\\Lambda$CDM) universe of side 2000 $h^{-1}$Mpc (3000 $h^{-1}$Mpc). Such large volumes allow many independent light-cones to be generated out to $z \\sim 1$. The light-cone output automatically accounts for clustering evolution with redshift. The principal uncertainty lies in how to create a `galaxy' distribution from the simulated mass distribution. We employ Lagrangian bias schemes similar to those of White et al.  (1987) and Cole et al.  (1998). Individual particles are tagged as galaxies with a probability which depends only on the smoothed {\\it initial} overdensity field in their neighbourhood. The parameters of these schemes are adjusted so that the present-day correlations of the simulated galaxies match observation. Many quasi-independent mock pencil-beam surveys can then be created adopting the geometry and the galaxy selection probability with distance of the BEKS surveys.  Our discussions focus primarily on the significance of the BEKS data in comparison with our CDM samples. We begin by following the methods used originally by BEKS, namely, redshift counts, pair separation distributions, and the one-dimensional power spectrum. Redshift counts are translated into a distribution in physical distance assuming the same cosmological parameters as BEKS. For the one-dimensional power spectra, the height of the highest peak is the most important statistic. Szalay et al. (1991) show that the statistical significance of the highest peak of the BEKS data is at $10^{-4}$ level, based on the formal probability for the peak height. This calculation was disputed by Kaiser\\&  Peacock (1991) because of the difficulty in estimating the appropriate noise level. We calculate relative peak-to-noise ratios of the highest peaks in the power spectra in identical ways for real and simulated data and so can compare the two without needing to resolve this issue. We also apply two additional statistical tests for regularity, the $\\Delta$ test of PG and a `supercluster' statistic designed by Dekel et al. (1992).  Our paper is organised as follows. In section 2 we present details of the $N$-body simulations from which our pencil-beam samples are drawn. In section 3 we explain our bias scheme. In section 4 we describe our mock pencil-beam surveys which mimic as closely as possible the actual observations of BEKS. Our main results for power spectra are given in section 5. Results are given in section 6 for the $\\Delta$ test, and in section 7 for the `supercluster' statistic. We present our conclusions in section 8. ",
        "conclusions": "By creating a number of mock pencil-beam surveys we have compared the apparent periodicity in two CDM model universes with that observed in the data of Broadhurst et al. (1990). The power spectrum analysis alone shows that the BEKS data are significantly more periodic than the models at about the half per cent level, while the PG $\\Delta$-test shows less significance, about 10 per cent for {\\bf t1} and 5 per cent for {\\bf L1} and {\\bf L2}. The supercluster statistic gives a two per cent probability of finding a structure as regular as the BEKS data. Restricting to a length scale $\\sim$100-150$h^{-1}$Mpc, however, the number of samples which show the kind of periodicity seen in the BEKS data is extremely small for each of these statistics. Overall no sample is more regular than the BEKS data for {\\it all} three statistics for a single period. The two popular CDM models we have studied here are apparently unsuccessful in reproducing the observed periodicity. From this result together with the fact that the statistical results appeared to be insensitive to the choice of the bias model, we conclude that CDM models conflict with the BEKS observation. Either the models need additional physics, or the data are a fluke or are somehow biased.  Various possible physical explanations have been proposed, such as coherent peculiar velocities (Hill, Steinhardt and Turner 1991) oscillations in the Hubble parameter (Morikawa 1991) or baryonic features in the power spectrum (Eisenstein \\etal 1998) but all of them seem to require either additional mechanisms with fine tunings beyond the standard theory or cosmological parameters significantly different from currently favoured values. Intriguingly, Dekel et al. (1992) demonstrated that built-in power on a large ($\\sim 100 h^{-1}$Mpc) length scale in the initial density fluctuation could indeed reproduce periodic features on a given scale, at least by some of the tests we have considered. If such excess power on large scales (hence still in the linear regime) exists, it will be detectable in the power spectra of the future 2dF and Sloan surveys.  Having found at least a few examples that are nearly as periodic as the BEKS data, we cannot rule out the possibility that the BEKS data (or the Galactic Pole direction) are a fluke. On the other hand, one should be aware of the complexity of the original observations --an incomplete compilation of a narrow and deep, and of a wide and shallow survey at each of the Galactic Poles. It is not clear whether such a combination constitutes a fair sample. No evidence for periodic structure on $\\sim$130$h^{-1}$Mpc has been found so far in two other deep redshift surveys, the ESO-Sculptor Survey (Bellanger and de Lapparent 1995) and the Caltech Faint Galaxy Redshift Survey (Cohen 1999). Follow-up observations to BEKS by Koo \\etal (1993) did {\\it not} show a strong regularity in two other directions, although around the Galactic Pole the regularity was found to be further strengthened. Our results give the {\\it a priori} probability for such apparent periodicity in CDM models. Several more deep surveys might suffice to judge whether the discrepancy with BEKS reflects a major inconsistency. The planned VIRMOS Deep Survey (Le F\\`{e}vre \\etal 1998, see also Guzzo 1999) will survey the range $0.3\\le z\\le 1$ and will provide, together with the large volume 2dF and Sloan surveys, much larger and more complete samples in the near future."
    },
    "0011/hep-ph0011097_arXiv.txt": {
        "abstract": "One of the possible consequences of the existence of extra degrees of freedom beyond the electroweak scale is the increase of neutrino-nucleon cross sections ($\\sigma_{\\nu N}$) beyond Standard Model predictions. At ultra-high energies this may allow the existence of neutrino-initiated extensive air showers. In this paper, we examine the most relevant observables of such showers. Our analysis indicates that the future Pierre Auger Observatory could be potentially powerful in probing models with large compact dimensions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011107_arXiv.txt": {
        "abstract": "Relativistic kinetic theory, based on the Grad method of moments as developed by Israel and Stewart, is used to model viscous and thermal dissipation in neutron star matter and determine an upper limit on the maximum mass of neutron stars. In the context of kinetic theory, the equation of state must satisfy a set of constraints in order for the equilibrium states of the fluid to be thermodynamically stable and for perturbations from equilibrium to propagate causally via hyperbolic equations. Application of these constraints to neutron star matter restricts the stiffness of the most incompressible equation of state compatible with causality to be softer than the maximally incompressible equation of state that results from requiring the adiabatic sound speed to not exceed the speed of light. Using three equations of state based on experimental nucleon-nucleon scattering data and properties of light nuclei up to twice normal nuclear energy density, and the kinetic theory maximally incompressible equation of state at higher density, an upper limit on the maximum mass of neutron stars averaging 2.64 solar masses is derived. ",
        "introduction": "The problem of the maximum possible mass of neutron stars has a long history. Baade and Zwicky in 1934 proposed a star composed of densely packed neutrons as the final state resulting from the supernova process \\cite{Baade34}. Oppenheimer and Volkoff in 1939 demonstrated a star composed of noninteracting neutrons is supported against gravitational collapse to a black hole by the Fermi degeneracy pressure only for stellar masses up to 0.72 solar masses ($M_\\odot$) \\cite{Oppenheimer39}. The discovery of the first pulsar in 1967 and the realization pulsars are highly magnetized rotating neutron stars triggered substantial theoretical work towards understanding the structure of neutron stars. Pulsar masses can be measured from observations of Doppler shifts of periodic emissions from neutron stars in binary systems. Some 20 radio pulsar masses have been accurately determined so far, with most masses clustering around 1.4 $M_\\odot$ \\cite{Thorsett99}. This is well beyond the value of 0.72 $M_\\odot$ computed for noninteracting neutrons, so the short-range repulsion of the strong nuclear force must make a substantial contribution to the pressure supporting neutron stars against gravitational collapse. There is also evidence for neutron stars with substantially larger masses than the radio pulsars due to accretion of matter from a binary companion. The mass of the x-ray pulsar in the Vela X-1 binary has been estimated to be 1.9 $M_\\odot$ \\cite{vanKerkwijk00}, and the x-ray burster Cygnus X-2 has been determined to be 1.8 $M_\\odot$ \\cite{Orosz00}. Quasiperiodic oscillations in the x-ray emissions from neutron stars accreting matter from low mass companions have been argued to imply a mass of up to 2.3 $M_\\odot$ for the neutron stars in these systems \\cite{Zhang97,Miller98}.  An accurate theoretical determination of the maximum possible mass of neutron stars is of practical interest for identifying as a black hole any compact object with a larger mass. The maximum mass value remains uncertain because the equation of state of neutron star matter is not well understood at the high density values found in neutron star interiors. Modern methods develop the equation of state by fitting experimental nucleon-nucleon scattering data and properties of light nuclei to two and three-body interaction potentials \\cite{Wiringa88,Akmal98}. Confidence in the results of these methods is high near normal nuclear density (energy density = 152 MeV/fm$^3$ = 2.7 X 10$^{14}$ g/cm$^3$, baryon density = 0.16 fm$^{-3}$) because, at low density, experimental verification is possible with existing facilities. The validity of extrapolating the experimentally verified low density equation of state to high density, or as an alternative theoretically modeling the physical processes thought to occur at high density,  is uncertain because of the lack of high density laboratory data. The sources of the uncertainties in the high density equation of state are incomplete knowledge of the three-nucleon interactions, the contributions of mesonic and other baryonic besides nucleonic degrees of freedom, and a possible quark deconfinement transition \\cite{Akmal98}.  In light of these uncertainties in the exact form of the high density equation of state, several authors have derived an upper limit on the maximum neutron star mass using only general restrictions on the equation of state. Oppenheimer and Volkoff suggested \\cite{Oppenheimer39}, and Rhoades and Ruffini rigorously proved \\cite{Rhoades74}, the maximum mass of a stable neutron star results when the stiffest equation of state compatible with thermodynamic stability and causality is used. Rhoades and Ruffini derived an upper limit on the maximum neutron star mass of 3.2 $M_\\odot$ by using the experimentally verified Harrison-Wheeler equation of state \\cite{Harrison65} up to 1.7 times normal nuclear energy density, and the stiffest equation of state compatible with causality at high density. They defined a casual equation of state to be one having an adiabatic sound speed less than the speed of light, so they took the most incompressible equation of state to have  the sound speed equal to light speed. Kalogera and Baym \\cite{Kalogera96} updated the result of Rhoades and Ruffini, using an equation of state developed by Wiringa, Fiks, and Fabrocini \\cite{Wiringa88} up to twice normal nuclear energy density. They derived an upper mass limit of 2.9 $M_\\odot$ using the same maximally incompressible equation of state as Rhoades and Ruffini above twice normal nuclear energy density.  The purpose of this paper is to apply insights from relativistic kinetic theory to the problem of the maximum possible mass of neutron stars. Kinetic theory used to model viscosity and heat conduction shows requiring the sound speed to be less than the speed of light is a necessary but not sufficient condition for causality. The equation of state must satisfy a set of constraining conditions for the fluid to have stable equilibrium states with perturbations from equilbrium that propagate causally via a system of hyperbolic equations \\cite{Hiscock83,Olson90}. If any one of these constraints is violated, there will be at least one fluid mode with a superluminal propagation speed. In particular, it is possible for the fluid to be acausal in a transverse mode or a different longitudinal mode even though the sound speed is below light speed. These constraints generally fix the stiffest possible equation of state to have an adiabatic sound speed significantly less than the speed of light. The main result of the present paper is application of these constraints to neutron star matter at high density softens the maximally incompressible equation of state and reduces the maximum neutron star mass significantly below the value found using a high density equation of state having the adiabatic sound speed equal to the speed of light. In the next section the thermodynamic constraint conditions from relativistic kinetic theory are reviewed. Then the maximally incompressible equation of state of neutron star matter is developed from these constraint conditions. Finally, neutron star models are constructed using this stiffest possible equation of state at high density, and the resulting maximum neutron star mass is determined. Gravitational units having $G=c=1$ are used. ",
        "conclusions": "The principal finding of this paper is stability and causality constraints resulting from applying kinetic theory to thermal and viscous dissipative processes in neutron star matter yields a softer maximally incompressible equation of state, and a lower maximum mass of neutron stars, than for the sound speed equal to light speed maximally incompressible equation of state. The reduction in the maximum mass averages from 2.93 $M_\\odot$ down to 2.64 $M_\\odot$ for the three low density equations of state considered if the low density equation of state is regarded as experimentally verified up to twice normal nuclear energy density.  The softening of the maximally incompressible equation of state is a generic prediction of kinetic theory, but the specific value of the resulting maximum neutron star mass derived here rests upon the accuracy of approximating the relaxation times and viscous-thermal couplings with their noninteracting degenerate Fermi gas forms, and the overall validity of dissipative relativistic kinetic theory for strongly interacting nuclear matter at high density. Future data from the Relativistic Heavy Ion Collider holds promise for experimental testing of the applicability of the kinetic theory approach to modeling high density nuclear matter.  Only static neutron star mass models have been constructed in this paper. Rotation can support larger mass stars against collapse, but only for stars rotating rapidly enough to be nearly shedding mass from the equator is the maximum mass value significantly changed from its nonrotating value. For these rapidly rotating stars the increase in the maximum mass is at most approximately 20\\% over the static value when undergoing uniform rotation \\cite{Cook92,Cook94}. Differential rotation is strongly damped \\cite{Hegyi77}. The remnant of binary neutron star coalescence may undergo a short period of dynamically stable differential rotation, and thus briefly support a much larger mass neutron star remnant against prompt black hole collapse than is possible for a static or uniformly rotating star \\cite{Baumgarter00}."
    },
    "0011/astro-ph0011331_arXiv.txt": {
        "abstract": "We have used the BIMA array to image the Galactic Center with a 19-pointing mosaic in HCN(1-0), HCO$^+$(1-0), and H 42$\\alpha$ emission with 5 km s$^{-1}$ velocity resolution and $13'' \\times 4''$ angular resolution.  The $5'$ field includes the circumnuclear ring (CND) and parts of the 20 and 50 km s$^{-1}$ clouds.  HCN(1-0) and HCO$^+$ trace the CND and nearby giant molecular clouds while the H 42$\\alpha$ emission traces the ionized gas in Sgr A West.  We find that the CND has a definite outer edge in HCN and HCO$^+$ emission at $\\sim45''$ radius and appears to be composed of two or three distinct streams of molecular gas rotating around the nucleus.  Outside the CND, HCN and HCO$^+$ trace dense clumps of high-velocity gas in addition to optically thick emission from the 20 and 50 km s$^{-1}$ clouds.  A molecular ridge of compressed gas and dust, traced in NH$_3$ emission and self-absorbed HCN and HCO$^+$, wraps around the eastern edge of Sgr A East.  Just inside this ridge are several arcs of gas which have been accelerated by the impact of Sgr A East with the 50 km s$^{-1}$ cloud.  HCN and HCO$^+$ emission trace the extension of the northern arm of Sgr A West which appears to be an independent stream of neutral and ionized gas and dust originating outside the CND. Broad line widths and OH maser emission mark the intersection of the northern arm and the CND.  Comparison to previous NH$_3$ and 1.2mm dust observations shows that HCN and HCO$^+$ preferentially trace the CND and are weaker tracers of the GMCs than NH$_3$ and dust.  We discuss possible scenarios for the emission mechanisms and environment at the Galactic center which could explain the differences in these images. ",
        "introduction": "A 2-5 pc circumnuclear ring (CND) of neutral gas and dust has been imaged at the Galactic Center in the far-infrared and radio \\citep{gen85,gus87,mar93}.  It has been proposed that the CND is part of an accreting disk of material surrounding the massive central source Sgr A$^*$ \\citep{gus87}, an apparent supermassive black hole \\citep{ghe98,gen97}.  Ionized gas streamers appear to be infalling towards the central mass, originating at the inner edge of the CND and converging at the location of the bar near Sgr A$^*$ \\citep{loe83}.  It is unclear whether the CND is a transitory or stable structure.  If it is stable, the question remains as to how the CND itself is maintained.  At a projected distance of 10 pc from the Galactic Center lie two giant molecular clouds (GMC), the 20 km s$^{-1}$ cloud (M-0.13-0.08) and the 50 km s$^{-1}$ cloud (M-0.02-0.07) \\citep{gus81} to the south and east of the CND respectively. \\citet{gus81b} argued that these clouds must lie more than 100 pc from the Galactic Center in order not to be tidally disrupted. Their interpretation of formaldehyde absorption data places the 50 km s$^{-1}$ cloud behind the nucleus and the 20 km s$^{-1}$ cloud in front. Various studies have investigated the interactions between these molecular clouds, the non-thermal radio source Sgr A East, the CND and the ionized gas streamers in the nucleus \\citep{hoe85, gen90, hoe91, ser92, den93, mar95, coi00}.  There are suggestions of molecular gas steamers from both of the GMCs located 10 pc away in projection that may be feeding the CND.  A southern streamer has been mapped in NH$_3$(1,1), (2,2) and (3,3) and sub-millimeter continuum \\citep{hoe91, den93, coi99} and emission in the same region has been seen in HCN(3--2) \\citep{mar95}.  An eastern streamer, lying at the edge of the Sgr A East SNR and protruding in towards the nuclear region, has been mapped in HCN(1--0) \\citep{hoe93} and HCN(4--3) and HCN(3--2) \\citep{mar95}.  This streamer has not been seen in NH$_3$ or sub-millimeter continuum.  More recently, \\citet{coi00} presented images of two streams of NH$_3$(1,1) and (2,2) emission parallel to the Galactic plane, which they interpret as evidence for infall from the 20 km s$^{-1}$ cloud onto the CND, and as evidence for interactions between Sgr A East and the 50 km s$^{-1}$ cloud, and a SNR to the south of Sgr A East (G359.92--0.09). They argue that both clouds are within 10 pc of the nucleus and that only the 20 km s$^{-1}$ cloud is interacting with the CND.    Many of these observations have a limited field-of-view and velocity coverage, mapping only the central $2-4'$ (5--10pc) of the Galaxy and not extending into the large molecular clouds, thereby unable to thoroughly address the question of possible interactions with the circumnuclear region.  This paper describes observations with the BIMA array\\footnote{The BIMA array is operated by the Berkeley-Illinois-Maryland Association under funding from the National Science Foundation.} of the central $5'$ in HCO$^+$, HCN(1-0), and H 42$\\alpha$ emission with 5 km s$^{-1}$ resolution, and in 87 GHz continuum emission.  The $5'$ field of view corresponds to a field roughly 12.5 pc at the distance of the Galactic Center (assuming $R_\\odot$ = 8.5 kpc). ",
        "conclusions": "i) We present images of HCN(1-0) and HCO$^+$ emission and absorption which trace dense molecular gas in the central 12 parsecs of the Galaxy.  The HCN(1-0) and HCO$^+$ emission primarily traces filamentary structures associated with the CND and nearby GMCs; low-level, extended emission from partially-resolved larger structures is also seen.  ii) The CND appears as a bright, well-defined ring of emission $20'' - 60''$ from the center, at an inclination of 60--70$^{\\circ}$.  Although the outer edge of the ring is not as sharply defined as the inner edge, the CND does not extend beyond $100''$ in HCN(1-0) and HCO$^+$ emission.  The CND does not appear to be an equilibrium structure;  rather, it consists of two or three separate streamers in rotation around  the nucleus.  The CND may exist long-term in a fluctuating sense, as gas streamers feed the inner parsec.  iii)  Outside the CND, HCN(1-0) and HCO$^+$ trace several filamentary structures which do not follow the rotation pattern of the $45''$ ring.  High-velocity emission is seen $90''$ east of Sgr A*, just inside the dense molecular ridge of shocked material seen at the intersection of Sgr A East and the 50 km sec$^{-1}$ cloud.  A stream of emission is also seen extending from the ionized northern arm, through the CND to high-velocity red- and blue-shifted material to the north. This material may be falling into the center on an orbit which intersects the CND streamers.  iv) Comparison of the HCN(1-0) and HCO$^+$ with NH$_3$ and dust emission shows that NH$_3$ emission is well correlated with self-absorbed HCN(1-0) and HCO$^+$ emission, and dust ridges in the nearby GMCs.  v) HCN(1-0) and HCO$^+$ emission appear to be enhanced in the CND relative to NH$_3$. Further observations of other spectral lines are required to distinguish between opacity, excitation and chemical explanations for these differences.   These observations show that the accretion of material into the Galactic Center region may be a chaotic and transient phenomenon.  From the spatial distribution of the nearby gas with respect to Sgr A East and other nearby supernova remnants, and from the overall kinematics and especially the isolated pockets of high-velocity gas, we have the impression that the nearby molecular clouds are being disrupted by the expanding shells associated with Sgr A East.  Infall of material appears to approach the CND from a number of different directions on different orbits.  Sgr A East itself may be driven by multiple outflow events associated with the central black hole or multiple supernovae.  We favor a model where the accretion process is highly non-steady and non-uniform at least on the timescale of the rotation period of the CND, 10$^{5}$ yrs.  In this scenario, the CND does not appear uniform because material continues to fall into the CND on short timescales.  Whether the further infall of material from the CND toward the central black hole is also transient and non-steady in nature is unclear from these studies. However, there is some evidence that material may approach the central region directly from outside of the CND via individual streamers.  {\\it Acknowledgments} This work was supported in part by NSF Grant AST-21795 to the University of California.  ALC is supported by an NSF Graduate Research Fellowship. We thank Robert Zylka for permission to use the 1.2 mm image.   \\newpage"
    },
    "0011/astro-ph0011041_arXiv.txt": {
        "abstract": "N-body + SPH simulations are used to study the evolution of dwarf irregular galaxies (dIrrs)  entering the dark matter halo of the Milky Way or M31 on plunging orbits. We propose a new dynamical mechanism driving  the evolution of gas rich, rotationally supported dIrrs, mostly found at the outskirts of the Local Group (LG), into gas free, pressure supported dwarf spheroidals (dSphs) or dwarf ellipticals (dEs), observed to cluster around the two giant spirals. The initial model galaxies are exponential disks embedded in massive dark matter halos and reproduce nearby dIrrs. Repeated tidal shocks at the pericenter of their orbit partially strip their halo and disk and trigger dynamical instabilities that dramatically reshape their stellar component. After only 2-3 orbits low surface brightness (LSB) dIrrs are transformed into dSphs, while high surface brightness (HSB) dIrrs evolve into dEs. This evolutionary mechanism naturally leads to the morphology-density relation observed for LG dwarfs. Dwarfs surrounded by very dense dark matter halos, like the archetypical dIrr GR8, are turned into Draco or Ursa Minor, the faintest and most dark matter dominated among LG dSphs. If disks include a gaseous component, this is both tidally stripped and consumed in periodic bursts of star formation. The resulting star formation histories are in good qualitative agreement with those derived using HST color-magnitude diagrams for local dSphs. ",
        "introduction": "Dwarf galaxies in the Local Group (LG) clearly obey a morphology-density relation.  Close to the Milky Way and M31 we find early-type dwarf galaxies, namely faint ($M_B > -14$) low surface brightness dwarf spheroidals (dSphs) and more luminous ($M_B > -17$), higher surface brightness dwarf ellipticals (dEs). All these galaxies are nearly devoid of gas, contain dark matter and mainly old stars and are supported by  velocity dispersion (Ferguson \\& Binggeli 1994, hereafter FB94; Mateo 1998, hereafter Ma98; Grebel 1999, hereafter Gr99; Van den Bergh 1999). Among them Draco and Ursa Minor have the highest dark matter densities ever measured (Lake 1990). On the outskirts of the LG we find similarly faint ($M_B>$ -18) and dark matter dominated dwarf irregular galaxies (dIrrs), that are gas rich, star-forming systems with disk-like kinematics (Ma98, Van den Bergh 1999, Gr99).   Previous attempts to explain the origin of dSphs in the LG have relied on gas dynamical processes to remove the gas in dIrrs.  Gas stripping may result either because of the pressure exerted by an external hot gaseous medium in the halo of the Milky Way (``ram pressure'') (Einasto et al. 1974) or because of internal strong supernovae winds (Dekel \\& Silk 1986).  However, ram pressure would require an external gas density that is several orders of magnitude higher than recently inferred for the Milky Way (Murali 2000) and supernovae winds cannot explain the existing morphology-density relation.  Moreover, such dissipative mechanisms would remove the gas but would not directly alter the structure and kinematics of the pre-existing stellar component. However, the light follows an exponential profile in both dSphs and dIrrs (Faber \\& Lin 1983; Irwin \\& Hatzdimitriou 1995; Ma98) and a positive correlation between surface brightness and luminosity is shown by both types of dwarfs (FB94), suggesting an evolutionary link between them. Is there a mechanism that can transform dwarf galaxies between morphological classes or must we accept the idea that dSphs are fundamentally different from dIrrs ?   Within rich galaxy clusters, fast fly-by encounters with the largest galaxies can transform a disk system into a spheroidal or S0 galaxy in just 3-4 Gyr (Moore et al. 1996, 1998).  If the halos of bright galaxies were scaled down versions of galaxy clusters then this ``galaxy harassment'' would be equally important within them. However, whereas rich clusters contain over thirty large ($L_*$) perturbing galaxies, the Milky Way and M31 have only a couple of satellites sufficiently massive to harass the other dwarf galaxies (Moore et al. 1999; Klypin et al. 1999) As a result, the rate for effective satellite-satellite fly-by encounters is less than one in every 10 Gyr (the LMC and the SMC being a notable exception).  Thus, we are left only with the repeated action of tidal forces from the primary galaxy as an evolutionary driver. These operate on the orbital timescale, which is of order of 3-4 Gyr in both clusters and galactic halos.  However, given the relatively low age of large, virialized clusters, galaxies have typically approached the cluster center only once by the present time, while  dSphs satellites have had sufficient time to complete several close tidal encounters with the Milky Way, as stellar ages imply that the latter was already in place  10 Gyr ago (Van den Bergh 1996).   In this paper we use  very high resolution N-Body + SPH simulations performed with the parallel binary treecode GASOLINE (Dikaiakos \\& Stadel 1996; Wadsley et al. 2000) to follow the evolution of small galaxies resembling dIrrs as they move on bound orbits in the tidal field of the massive dark matter halo of the Milky Way. ",
        "conclusions": "Figure 5 summarizes the main observable properties of the simulated satellites projecting them on the Fundamental Plane (FB94). The remnants of LSB satellites resemble dSphs like Fornax or Sagittarius ($-14 < M_B < -11$), while HSBs transform in the bright dEs ($M_{B} > -17$), having a final central surface brightness higher than that of observed dIrrs with the same luminosities and therefore matching another observational constraint (FB94;Ma98). The total (including dark matter) final mass-to-light  ratios are in the range $6-20$. Remarkably, our model can reproduce the properties of even the most extreme dSphs, Draco and Ursa Minor.  In fact,  as the dark matter halo of GR8 is barely affected by tides (Figure 1), the remnant ($M_B = -7.5$) has a final mass-to-light ratio still $\\sim 50$ and the central dark matter density is still around $0.3 M_{\\odot}$ pc$^{-3}$, matching the structural parameters inferred for Draco and Ursa Minor (Ma98).  \\medskip \\epsfxsize=8truecm \\epsfbox{fundpnew.ps} \\figcaption[fundpnew.ps]{\\label{fig:asymptotic} \\small{Fundamental Plane (FP) for all the remnants as as in Ferguson \\& Binggeli (1994) (bottom). In the FP plot, $\\mu$ and  $\\sigma$ are, respectively, the average surface brightness and velocity dispersion measured inside $R_e$ and $M_B$ is measured at the Holmberg radius. The dashed line is a fit to the distribution of local dSphs, while the solid line refers to elliptical galaxies within the Virgo cluster (Dressler et al. 1987). Open and filled symbols represent, respectively, the initial and final state of HSBs (squares) and LSBs (circles). The initial and final state of the GR8 model are indicated by the open triangle and ``star''.}}     ``Tidal stirring'' naturally leads to the spatial segregation of dIrrs versus dSph as its effectiveness depends strongly on the distance from the primary.  How important is our assumption of a massive and extended dark matter halo surrounding the Milky Way ? When we adopt a ``minimal'' dark halo truncated at 50 kpc (with mass $5 \\times 10^{11} M_{\\odot}$; Little \\& Tremaine 1987), tides are too weak and the final remnants are still rotationally flattened ($v/\\sigma > 1$).  Instead, within a halo as massive and extended as implied by theories of galaxy formation (Peebles 1989) , our dIrrs models transform into dSphs even on orbits with apocenters larger than 200 kpc, explaining the origin of even the farthest dSphs as Leo I and Leo II.   Though rather speculative at this stage, it is tempting to relate HSB satellites observed during the strong bursting phase to the population of blue compact dwarfs identified by Guzm\\'an et al. (1997) at intermediate redshift. Redshift surveys will establish if bursting dwarfs have nearby massive companions.  Extended tidal streams of stars originate from our simulated dwarfs (Fig. 2) with a maximum surface brightness of just 30 mag arcsec$^{-2}$ (B band).  Spectroscopic evidence for stellar streams from the dSph Carina has been recently claimed (Majewski et al. 2000). Future astrometric missions, like SIM and GAIA (Gilmore et al. 1998; Helmi et % should reveal such faint features and will also carry out high-quality measurements of proper motions for many satellites of the Milky Way, thus providing a test for the orbital configurations used in this model.  Our model successfully explains the origin of dSphs once all observational constraints are taken into account: they evolved from dIrrs that entered the halo of the Milky Way or M31 several Gyr ago moving on plunging orbits and  suffered stirring by the tidal field of the large spirals.  \\vskip 16 pt  The authors thank G.Bothun for stimulating discussions. Simulations were carried out at the CINECA (Bologna) and ARSC (Fairbanks) supercomputing centers."
    },
    "0011/astro-ph0011277_arXiv.txt": {
        "abstract": "\\noindent From theoretical predictions, the introduction of 8--m class telescopes permits one to extend Surface Brightness Fluctuations measurements from the ground to $\\approx$7000 km/s, with a precision that is comparable to current space--based measurements. We have measured I--band SBF in IC 4296 in Abell 3565 $cz \\sim 3630$ km/s with the ESO Very Large Telescope. Adopting the Tonry et al. 2000 calibration for I--band SBF we determined a distance modulus of $(\\overline I_{o,k} - \\overline M_{I}) = 33.44 \\pm 0.17$ corresponding to a galaxy distance of 49 $\\pm$ 4 Mpc. This result is consistent with the HST observation from Lauer et al. 1998: $(\\overline I_{F814} - \\overline M_{F814}) = 33.47 \\pm 0.13$. ",
        "introduction": "\\noindent I--band Surface Brightness Fluctuations (SBF) have been successfully used  to measure early-type galaxy distances up to 4000 km/s from the ground and up to 7000 km/s with the Hubble Space Telescope  \\cite{sod95,sod96,aj97,ton97,tho97,lau98,pah99,bla99a,ton00}. This method was introduced by Tonry \\& Schneider 1988 (a recent review has been given by Blakeslee et al. 1999a) and is based on a simple concept. 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 like this variance normalized to the mean flux of the galaxy in each pixel \\cite{ts88}. The stars that contribute the most to the SBF are then the brightest stars in the intrinsic luminosity function, typically red giants in old stellar populations.  The absolute magnitude of the fluctuation is not a constant, but depends on the age and metallicity of the stellar population. As such it is a strong function of not only the photometric band in which the observations are carried out, but also of the color of the galaxy under study. Tonry et al. (1997) and Tonry et al. (2000) 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) color. The absolute calibration is then obtained using a set of 5 galaxies with independent Cepheid distance. The I--band SBF absolute magnitude calibration can be written as: \\begin{equation} \\overline M_I=-1.74+4.5\\ [(V-I)_0-1.15]. \\label{eq:Ical} \\end{equation} \\noindent 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. The Cepheid distances were part of a larger dataset of 18 distances to galaxies within the Fornax and Virgo clusters observed by the HST Key Project on the Extragalactic Distance Scale. The two calibrations are consistent within the errors.  \\small{  \\begin{table*}[!htf]  \\begin{flushleft} \\begin{tabular} {|c|c|c|c|c|c|c|c|c|} \\hline Name&RA &DEC &l&b&Type&$m_V$&$V-I$&$V_{LG} $ \\\\ \\hline &(2000)&(2000)&deg&deg&&mag&mag&km/s \\\\ \\hline IC 4296&13 36 39.46 & -33 57 59.8 &313.54 &+27.97&E&10.57&1.24&3630 \\\\ \\hline  \\end{tabular}  \\end{flushleft}   \\caption{General properties of IC 4296, from the web archive Simbad (http://simbad.u-strasbg.fr), and Lauer et al. 1998} \\label{tab-vlt} \\end{table*} } \\normalsize \\noindent The potential of these measurements extends to the measurements of cosmological parameters and the study of the dynamics of the universe. SBF distances have been used in measuring the Hubble constant in the HST Key project on the Extragalactic Distance Scale (Ferrarese et al. 2000a, and references therein). Tonry and his collaborators \\cite{ton00} have derived peculiar velocities from their ground based sample up to 4000~km/s and calculated bulk flows in this region. By comparing these peculiar velocities with infrared and optical surveys, Blakeslee et al. 1999a have been able to constrain the value of the Hubble constant and the parameter density (luminous plus dark matter) $\\beta=\\Omega^{0.6}/b$, where $b$ is the linear bias \\cite{bla99b}.  \\noindent However, the current I--band sample is limited to 4000~km/s. The introduction of large telescopes, such as the Very Large Telescope (VLT), have opened new opportunities to extend the current sample from the ground. With the optical instrument FORS1 and FORS2 on VLT, from a theoretical point of view,  it will be possible to extend I--band SBF with high signal--to--noise and good understanding of external source contribution up to $\\approx$ 7000~km/s \\cite{mei00}. This implies that ground--based  measurements can be used to probe the same volume as HST measurements.  \\noindent To illustrate the potential of I--band SBF measurements with the VLT, we have observed IC 4296 in Abell 3565.  This galaxy  is part of the brightest ellipticals in the Abell clusters observed in the Lauer \\& Postman 1992 sample. An I--band SBF distance measurement of this galaxy, based on HST observations, has already been reported by Lauer et al. (1998). We compare our SBF measurements with Lauer et al. (1998) and obtain similar results and measurement precision, being our errors and external source contribution estimation comparable with HST capabilities. ",
        "conclusions": "\\noindent  I--band SBF have been observed in IC 4296 with the VLT. A distance modulus of $(\\overline I_{o,k} - \\overline M_I) = 33.44 \\pm 0.17$ was measured, from which a galaxy distance $49 \\pm 4$ Mpc is derived.  \\noindent This result confirm the potential of 8--m class telescopes in these kind of measurements and suggest that future SBF observations from the ground can reach the same distance that until now were only reachable by HST."
    },
    "0011/astro-ph0011457_arXiv.txt": {
        "abstract": "s{ The  evolution of the X-ray luminosity function of clusters of galaxies has been measured to z=0.9 using over 150 X-ray selected clusters discovered in the WARPS survey. We find no evidence for evolution of the luminosity function at any luminosity or redshift. The observations constrain the evolution of the space density of moderate luminosity clusters to be very small, and much less than predicted by most models of the growth of structure with $\\Omega_m$=1. All the current X-ray surveys agree on this result. Several notable luminous clusters at z$>$0.8 have been found, including one cluster which is more  luminous (and is probably more massive) than the well known MS1054 cluster. } ",
        "introduction": "The evolution of the space density of clusters of galaxies is a measurement sensitive to the physical and cosmological parameters of models of structure formation. We describe a deep X-ray survey of clusters of galaxies (the Wide Angle ROSAT Pointed Survey or WARPS), and use it to measure the evolution of the X-ray luminosity function (XLF) of clusters of galaxies (Scharf et al 1997, Jones et al 1998, Ebeling et al 2000a, Fairley et al 2000). ",
        "conclusions": "\\subsection{Comparison with results of X-ray cluster surveys}  Several recent deep X-ray clusters surveys are described in these proceedings and elsewhere. In the regime where these surveys have good statistical accuracy (ie. moderate X-ray luminosities $\\sim$10$^{44}$ erg s$^{-1}$) there is excellent agreement that {\\it no evolution of the XLF is observed to z$\\approx$0.8}. Five surveys agree on this point (EMSS, RDCS, SHARC, CfA, WARPS) and the only disagreement (the RIXOS survey of Castander et al 1995) can be understood in terms of the RIXOS source detection algorithm.  At the higher X-ray luminosities of the most massive clusters (L$_X>$5x10$^{44}$ erg s$^{-1}$), there is some disagreement between the results of different surveys as to the degree of evolution found at z$>$0.3 (if any). This may partly be due to the small numbers of high luminosity clusters in any one survey. The range of evolution found is not large: from none to negative evolution of a factor $\\approx$3.  The WARPS and RDCS (Borgani et al, these proceedings) are both consistent with no evolution of the XLF up to z=1. WARPS is not sensitive to clusters at z$>$1 because our optical followup does not extend into the NIR.  \\subsection{Revisiting the EMSS survey}  In an effort to understand the EMSS results of Gioia et al (1990) and Henry et al (1992), who found negative evolution at high luminosities but relatively low redshifts (z$\\approx$0.33), we have remeasured the X-ray luminosities of the 11 EMSS clusters at z$>$0.3 for which deep ROSAT PSPC data exists, extending the work started by Jones et al (1998). If the X-ray luminosities of EMSS z$>$0.3 clusters have been significantly underestimated by a factor of $\\approx$2, then the EMSS XLF will move toward a no-evolution result.  We use large (3 Mpc radius, H$_0$=50) apertures, removing point sources and using ASCA temperatures where necessary. We find that the mean ratio $L_{X,PSPC}/L_{X,EMSS}\\approx1.6$ and that the ratio is correlated with the core radius measured from the PSPC images. This suggests that the assumption of a constant core radius of 250 kpc in the EMSS (a reasonable assumption, given the data available at the time) has led to an underestimate of the luminosities of clusters with larger core radii.  Henry (2000) has investigated EMSS cluster luminosities and notes that the original EMSS extrapolation of the surface brightness profile to give the total flux (a mean correction factor of $\\approx$2.5) did not take into account the Einstein IPC psf. Including the effect of the psf revises the EMSS luminosities upwards by a factor of 1.37, explaining a major part of the discrepancy with the ROSAT luminosities.  However, Henry (2000) also notes that luminosities measured with ASCA are only 17\\% higher than EMSS luminosities.   \\subsection{Future work} Future work will concentrate on the handful of clusters in the WARPS survey at z$\\approx$0.8, on measurements of the evolution of the cluster X-ray temperature function using Chandra and XMM, and studies of galaxy evolution in X-ray selected clusters. We are also studying the WARPS sample of `fossil' groups of galaxies, to help understand the formation of luminous elliptical galaxies.   New wide area X-ray surveys, designed to detect the rare high luminosity, massive clusters in large numbers at high redshifts (z$>$0.5) are planned (see Lumb \\& Jones 2000 and Ebeling et al 2000c). It is the most massive clusters which have the greatest leverage to constrain cosmological parameters."
    },
    "0011/astro-ph0011396_arXiv.txt": {
        "abstract": "It is shown that, under sufficiently intense OB-star illumination of a stationary photoexcitation front (PDR), nonlinear H$_{2}$ photoexcitation processes comprising driven resonant two-photon transitions between X-state quantum levels, with VUV continuum light from the star supplying both driving fields, largely determine the photonic pathways of H$_{2}$ molecules in the PDR close to the ionization front. Specifically, for a flux of $% \\sim $4 x 10$^{5}$ Habing fields incident upon a PDR, the total rate at which an H$_{2}$ molecule is nonlinearly photoexcited out of any X-state quantum level is calculated to be roughly 100 times greater than the total rate at which it is linearly photoexcited out of the same level. In strongly excited PDRs, the populations in almost all of the $\\sim$300 bound quantum levels of the X state will be maintained approximately equal via a few myriads of interconnecting two-photon steps. The remarkable importance of two-photon transitions in H$_{2}$ photoexcitation in strongly irradiated PDRs derives from the exceptionally narrow Raman linewidth ($\\Gamma \\sim $10$^{-6}$ $\\sec ^{-1}$) that characterizes all two-photon transitions between bound H$_{2}$ X-state quantum levels. ",
        "introduction": "In existing theoretical models of stationary photoexcitation fronts (PDRs) \\setlength{\\baselineskip}{0.5\\normalbaselineskip} $[$see Draine \\& Bertoldi (1996) for a comprehensive list of references$]$, a one-dimensional geometry is usually considered, with light from an OB-type illuminating star assumed incident upon the PDR from one direction. In all published PDR models, only linear H$_{2}$ photoexcitation processes have been considered. The consensus among astrophysicists is that these models adequately describe the main photonic events that occur in H$_{2}$% -containing clouds separated by at least 10 pc from OB-type illuminating stars. In the present Letter, we draw attention to the fact that, in very strongly irradiated PDRs (for example, those with cloud-to-star distances $% \\leq $0.1 pc), nonlinear photoexcitation of H$_{2}$ can evidently become the dominant photonic process determining the structure of the PDR. ",
        "conclusions": ""
    },
    "0011/astro-ph0011169_arXiv.txt": {
        "abstract": "We report the presence of a 106-day cycle in the radio variability of Sgr A*  based on an analysis of data observed with the Very Large Array (VLA) over the past 20 years. The pulsed signal is most clearly seen at 1.3 cm with a ratio of cycle frequency to frequency width $f/\\Delta f= 2.2\\pm0.3$.  The periodic signal is also clearly observed at 2 cm. At 3.6 cm the detection of a periodic signal is marginal. No significant periodicity is detected at both 6 and 20 cm.  Since the sampling function is irregular we performed a number of tests to insure that the observed periodicity is not the result of noise. Similar results were found for a maximum entropy method and periodogram with CLEAN method.  The probability of false detection for several different noise distributions is less than 5\\% based on Monte Carlo tests. The radio properties of the pulsed component at 1.3 cm are spectral index $\\alpha\\sim1.0\\pm 0.1$ (for $S\\propto\\nu^\\alpha$), amplitude $\\Delta S=0.42\\pm 0.04 {\\rm\\ Jy}$ and characteristic time scale $\\Delta t_{FWHM}\\approx 25\\pm5$ days. The lack of VLBI detection of a secondary component suggests that the variability occurs within Sgr A* on a scale of $\\sim5$ AU, suggesting an instability of the accretion disk. ",
        "introduction": "The compact radio source Sagittarius A* is suggested to be associated with a supermassive black hole  at the Galactic Center (Eckart \\& Genzel \\etal 1997, Ghez \\etal 1998,  Backer \\& Sramek 1999, Reid \\etal 1999). The flux density variability of Sgr A* has been puzzling since the discovery of this intriguing radio compact source at the center of the Galaxy in 1974 (Brown \\& Lo 1982; Zhao \\etal 1989). We monitored Sgr A* with the VLA during the period of 1990 to 1993. Fluctuations in flux density suggested that the amplitude of variation increased towards short wavelengths and that the rate of outbursts is about three per year (Zhao \\etal 1992; Zhao \\& Goss 1993). Large amplitude fluctuations in the flux density have been observed at millimeter wavelengths (Backer \\& Wright 1993; Tsuboi \\etal 1999). Based on the radio monitoring data obtained with the 3.5 km Green Bank Interferometer at 11 and 3.6 cm, Falcke (1999) reported that at both wavelengths a characteristic time scale of 50 to 200 days is observed while the structure function of 11 cm data suggests a quasi-periodic variation with a period of 57 days. ",
        "conclusions": "If the 106 day cycle is related to an orbital emitting object around the massive black hole, its distance to the central mass would be 1200$\\times$ R$_g$ (60 AU), where R$_g = 7.5\\times 10^{11}$ cm for 2.5$\\times10^6$ M$_\\odot$. At a distance of 8 kpc, 60 AU corresponds to 8 mas.  A 200 mJy compact object separated by 8 mas from Sgr A* is easily detected with the VLBA at wavelengths shorter than 3.6 cm. There is no evidence for a companion source in images of Sgr A* at any wavelength (e.g., Bower \\& Backer 1998). We can also demonstrate the unlikelihood of the specific case of an eccentric binary pair in which a wind from the secondary eclipses the primary as in the case for the pulsar PSR 1259-63 (Johnson \\etal 1992). The opacity of free-free absorption $\\tau_{f-f} \\sim \\lambda^{2.1}$ suggests that the periodic variability due to an orbiting body would be enhanced at longer wavelengths. At $\\lambda \\ge$ 6 cm, Sgr A* would be completely absorbed in the eclipsed phase if $\\tau\\sim 0.5$ as inferred from the fraction of 40\\% in the flux density variation at 1.3 cm. Thus, the periodic variability is more likely intrinsic to Sgr A* and probably occurs within a few hundred R$_g$ of the massive black hole.   Quasiperiodic variability in radio flux density can be produced through jets or collimated flows. With a wealth of data observed at wavelengths between radio to X-ray, the microquasar GRS 1915+105 provides an excellent case showing two distinct states ( ``plateau'' and ``flare'') of the accretion process in a stellar mass black hole, (Mirabel \\& Rodriguez 1999; Dhawan \\etal 2000). The current state of Sgr A* is similar to the ``plateau'' state of GRS 1915+105 in terms of flat radio spectrum, compact source size (AU scale) and periodic variability in radio flux density. This state can be contrasted to  the ``flare'' state  which is characterized by optically thin ejecta feeding large scale jets (500 AU). The radio variability as correlated with the soft X-ray cycle in the ``plateau'' state of GRS 1915+105 suggests that the radio oscillations correspond to  the thermal-viscous instability of the accretion disk (Dhawan \\etal 2000). The analog of the radio fluctuations in the two sources suggests that the periodic flux density variations in Sgr A* are also related to an instability of the accretion disk. However, the X-ray luminosity ($\\sim2\\times10^6$ L$_\\odot$) of GRS 1915+105 is 20 times greater than the Eddington limit for  a 3 M$_\\odot$ black hole (Mirabel \\& Rodriguez 1999) and the radio jets are produced by the overwhelming radiation pressure. On the other hand, the low X-ray luminosity ($<$ 100 L$_\\odot$, $\\sim$10 orders in magnitude below the Eddington limit for a 2.5$\\times10^6$ M$_\\odot$ object) of Sgr A* indicates that the gravity   far exceeds the radiation pressure. Because of the strong gravity and weak radiation pressure, the consequence of the gas dynamics in Sgr A* on the AU scales would be different from  GRS 1915+105. In fact, the time variability and the limit  on the intrinsic source size  ($<$0.5 mas or 100 R$_g$ or 5 AU) of Sgr A* from the 7 mm observations (Lo \\etal 1998 and Bower and Backer 1998) suggest that any variability in the jet occurs in a region where the gravitational field of the black hole dominates. Any collimations of jets or outflows related to the observed radio variability appear to be  disrupted within Sgr A* on a scale of $\\sim$5 AU. A model consisting of  a jet nozzle ({\\it e.g.} Falcke 1996) seems useful to study if a self-consistent dynamic theory for the disk instability can be constructed.   A convection process is now considered in the advection dominated accretion flow (ADAF)  that could well provide a reasonable dynamic model for the accretion disk and possible outflows (Narayan \\etal 2000; Quataert \\& Gruzinov 2000; Igumenshchev \\& Abramowicz, 1999; Stone \\etal, 1999). In this theory, hot dense bubbles are produced in the inner part of a low-viscosity disk through convection caused by thermal instability. The most attractive result from the convective-ADAF theory is that quasiperiodic production of convective bubbles  has been observed in numerical simulations, although the observed period and the small cycle frequency width have  not been  predicted in detail from the theory (Igumenshchev \\& Abramowicz, 1999). The current results appear to favor the convective-ADAF model although we can not rule out the possibilities of an orbiting companion object that triggers periodic flares from a jet nozzle."
    },
    "0011/astro-ph0011443_arXiv.txt": {
        "abstract": "While all but one Gamma-Ray Bursts observed in the X-ray band showed an X-ray afterglow, about 60 per cent of them have not been detected in the optical band. We demonstrate that in many cases this is not due to adverse observing conditions, or delay in performing the observations.  We also show that the optically non-detected afterglows are not affected by particularly large Galactic absorbing columns, since its distribution is similar for both the detected and non-detected burst subclasses.  We then investigate the hypothesis that the failure of detecting the optical afterglow is due to absorption at the source location.  We find that this is a marginally viable interpretation, but only if the X-ray burst and afterglow emission and the possible optical/UV flash do not destroy the dust responsible for absorption in the optical band.  If dust is efficiently destroyed, we are led to conclude that bursts with no detected optical afterglow are intrinsically different.  Prompt infrared observations are the key to solve this issue. ",
        "introduction": "The standard external shock synchrotron model (Meszaros \\& Rees 1997; Sari et al. 1998) has been very successful in describing the properties of observed optical afterglows (Wijers et al. 1997; Galama et al. 1998d, Covino et al. 1999).  However, for more than half of the afterglows observed in the optical band we did not detect any emission.  We will call these Failed Optical Afterglow (FOA) gamma-ray bursts.  In all the \\grb~error boxes promptly followed by narrow field X-ray instruments, an X-ray transient (afterglow) has been detected (with the only exception of GRB~990217), while only for $\\sim 40$ per cent of them optical observations have revealed an afterglow at optical wavelengths.  This despite the rough similarity of the X-ray afterglow fluxes and the prompt reaction of optical telescopes. Paczynski (1998) ascribes this failed detection to dust extinction pointing out how this interpretation requires the association of bursts with star forming regions.  If this is the case, infrared observations should be better suited for the hunt of afterglows, where the extinction plays a reduced role.  For this reason, the IR follow-up of GRBs has recently become quite common, and some afterglows (GRB~990705, Masetti et al. 2000; and GRB~000418, Klose et al. 2000a) have been detected in the infrared before being confirmed at optical wavelengths.  Yet, we still miss the detection of an IR afterglow without an optical counterpart: such a detection would confirm the role of dust in FOAs.  Adding confusion to this picture, observations of extinction in X-ray spectra seem to reveal a very low gas to dust ratio (Vreeswijk et al. 1999, Galama \\& Wijers 2000) which, if common in all GRB environments, would strongly limit the role of dust extinction in the absorption of afterglows in the optical band and, even more, in the NIR.  In this paper we show that upper limits derived for FOAs are indeed not consistent with an ``average afterglow'', contrary to what recently claimed by Galama \\& Wijers (2000).  We then analyze the properties of detected and non detected optical afterglows in order to check whether the absorption commonly seen in star forming regions can explain the large fraction of FOAs.    \\begin{table*} \\begin{tabular}{lllllllllll} \\hline GRB  &$F_{X,\\rm NFI}$  &$\\Delta t_x$ &$\\delta_x$ &Ref  &$R$   &$\\Delta t_R$ &$\\delta_R$ &Ref &$z$ &Ref \\\\ &$10^{-13}$cgs  & h           &           &       &      &h            &           &    &    &      \\\\ \\hline \\hline 970228.12362 &28$\\pm$4    &8       &1.32       &Co97 &21.5$\\pm$0.3   &16.5  &1.73$\\pm$0.12 &Ma98,Ga00 &0.695   &Bl98  \\\\ 970508.904   &7$\\pm$0.7   &6       &1.1        &Am98 &19.77$\\pm$0.1   &52   &1.2$^a$       &Pe98,Ga98a &0.835  &Me97   \\\\ 971214.97272 &4$\\pm$0.4   &6.7     &0.9        &An97 &22.06$\\pm$0.06 &13    &1.20$\\pm$0.02 &Di98      &3.418   &Ku98   \\\\ 980326.88812 &NP          &---     &---        &---  &21.25$\\pm$0.03 &11    &2             &Bl99      &---     &---        \\\\ 980329.1559  &7.8$\\pm$0.9 &7       &1.35    &Za98a,b &21.2$\\pm$0.3$^b$ &17  &1.3$\\pm$0.1   &Re99      &---     &---    \\\\ 980425.90915$^c$ &4$\\pm$0.6 &10    &0.2       &Pi00a &15.7$\\pm$0.1   &59.8  &---           &Ga98b     &0.0085  &Ti98    \\\\ 980519.51403 &1.4$\\pm$0.3 &9.7     &1.8        &Ni99 &20.4$\\pm$0.1   &15.5  &2.05$\\pm$0.07 &Ha99      &        &   \\\\ 980613.20215 &1.1$\\pm$0.3 &9       &0.8        &Co99a &22.9$\\pm$0.2   &16.3  &1             &Hj98,Dj98a &1.0964 &Dj98b    \\\\ 990123.40780 &110         &5.8     &1.35  &He99a &18.26$\\pm$0.04$^d$ &3.8   &1.12$\\pm$0.03 &Od99,Ga99 &1.6004  &Ku99    \\\\ 990510.36743 &14.7$\\pm$1.8 &8      &1.4        &Ku00 &17.54$\\pm$0.02 &3.5   &0.82$\\pm$0.02 &Hr99      &1.619   &Vr99a   \\\\ 990705.66765 &1.9         &11      &1.6        &Am00 &18.7$\\pm$0.05$^e$ &5.5 &1.68$\\pm$0.10 &Ma00    &---     &---   \\\\ 990712.69655 &NP          &---     &---        &---  &19.4$\\pm$0.1   &4.16  &0.97$\\pm$0.02 &Sa00a     &0.4331  &Vr00 \\\\ 001011.66308 &NP          &---     &---        &---  &20.6$\\pm$0.1   &8.4   &1.4           &Go00      &---     &---    \\\\ \\hline 980703.182468 &7.5        &22    &1.3$\\pm$0.25 &Ga98c &21.00$\\pm$0.09 &22.6 &1.39$\\pm$0.3  &Ca99a      &0.9662  &Dj98c \\\\ 990308.21883  &---        &---     &---        &---  &18.14$\\pm$0.05 &3.34  &1.2$\\pm$0.1   &Sc99      &---     &---  \\\\ 991208.192269 &---        &---     &---        &---  &18.7$\\pm$0.1   &49.9  &2.15          &Je99a      &0.7055  &Di99 \\\\ 991216.671544 &1240$\\pm$40 &4.03   &1.64       &Ta99 &18.49$\\pm$0.05 &10.8  &1.22$\\pm$0.04$^f$ &Ha00  &1.02    &Dj99 \\\\ 000131.62446  &---        &---     &---        &---  &23.26$\\pm$0.04 &84.3  &2.25$\\pm$0.19 &An00      &4.50    &An00 \\\\ 000301.41084  &---        &---     &---        &---  &20.42$\\pm$0.06 &36.5  &1.18$\\pm$0.14 &Sa00b     &2.0335  &Ca00a \\\\ 000418.41921  &---        &---     &---        &---  &21.63$\\pm$0.04 &59.3  &0.86$\\pm$0.06 &Kl00      &1.1854  &Bl00 \\\\ 000630.02145  &---        &---     &---        &---  &23.04$\\pm$0.08 &21.6  &1.1$\\pm$0.3   &Je00      &---     &---  \\\\ 000911.30237  &---        &---     &---        &---  &20.26$\\pm$0.17 &34.3  &1.5$\\pm$0.14  &Pr00a,La00 &---     &---  \\\\ 000926.99274  &2.1$\\pm$0.6 &54.2   &4.3$\\pm$1.0 &Pi00b &19.37$\\pm$0.02 &20.7 &1.36$\\pm$0.11 &Sa00c,Fy00a &2.066 &Fy00b\\\\ 001007.20749  & ---       &---     &---        &---  &20.3           &83    &0$^g$         &Ca00b,Pr00b &---    &--- \\\\ \\hline \\hline \\end{tabular} \\begin{flushleft} Notes: $\\delta_\\nu$ is defined by $F_\\nu (t) \\propto t^{-\\delta_\\nu}$. X-ray fluxes in the 2--10 keV band. NP=repointing of \\sax~not possible. $^a$: for $t>2$ days; earliest detection at 3.1 hours: $R=21.1\\pm0.1$. $^b$: $R$ mag derived from $I=20.8\\pm0.3$; $R$=23.6$\\pm$0.2 after 20 hours. $^c$: = SN 1998bw, not used in the analysis. $^d$: Converted from Gunn $r$-mag. $^e$: $R$ mag derived from $H=16.57\\pm0.05$. $^f$: for t$\\le$1.2 days. $^g$: for t$\\le$3.5 days, $\\delta_R\\sim1.4$ after.\\\\ Am98: Amati et al., 1998; Am00: Amati et al., 2000 An97: Antonelli et al., 1997; An00: Andersen et al., 2000; Bl98: Bloom et al., 1998; Bl99: Bloom et al., 1999; Bl00: Bloom et al., 2000; Ca99a: Castro-Tirado et al., 1999a; Ca00a: Castro et al., 2000a; Ca00b: Castro et al., 2000b; Co97: Costa et al., 1997; Co99a: Costa et al., 1999a; Di98: Diercks et al., 1998; Dj98a: Djorgovski et al., 1998a; Dj98b: Djorgovski et al., 1998b; Dj98c: Djorgovski et al., 1998c; Dj99: Djorgovski et al., 1999; Do99: Dodonov et al., 1999; Fy00a: Fynbo et al., 2000a; Fy00b: Fynbo et al., 2000b; Ga98a: Galama et al., 1998a; Ga98b: Galama et al., 1998b; Ga98c: Galama et al., 1998c; Ga99: Galama et al., 1999; Ga00: Galama et al., 2000; Go00: Gorosabel et al., 2000; Ha99: Halpern et al., 1999; He99a: Heise et al., 1999; Hj98: Hjorth et al., 1998; Hr99: Harrison et al., 1999; Je99a: Jensen et al., 1999a; Je00: Jensen et al., 2000; Kl00: Klose et al., 2000a; Ku98: Kulkarni et al., 1998; Ku99: Kulkarni et al., 1999; Ku00: Kuulkers et al., 2000; La00: Lazzati et al., 2000; Ma98: Masetti et al., 1998; Ma00: Masetti et al., 2000; Me97: Metzger et al., 1997; Ni99a: Nicastro et al., 1999a; Od99: Odewahn et al., 1999; Pe98: Pedersen et al., 1998; Pi00a: Pian et al., 2000; Pi00b: Piro et al., 2000; Pr00a: Price et al., 2000a; Pr00b: Price et al., 2000b; Re99: Reichart et al., 1999; Sa00a: Sahu et al., 2000; Sa00b: Sagar et al., 2000; Sa00c: Sagar et al., 2001; Sc99: Schaefer et al., 1999; Ta99: Takeshima et al., 1999; Ti98: Tinney et al., 1998; Vr99a: Vreeswijk et al., 1999a; Vr00: Vreeswijk et al., 2001; Za98a: In't Zand et al., 1998a; Za98b: In't Zand et al., 1998b; \\end{flushleft} \\caption{{Properties of the bursts with associated optical transient. The first 13 bursts have been observed by the \\sax-GRBM/WFC, while the remaining bursts (below the horizontal line) have been discovered by other instruments (see text).} \\label{tab:uno}} \\end{table*}    \\begin{table*} \\begin{tabular}{llllllll} \\hline GRB &$F_{X,\\rm NFI}$ &$\\Delta t_x$ &$\\delta_x$ &Ref &$R$ &$\\Delta t_R$ &Ref  \\\\ &$10^{-13}$cgs    & h          &           &    &    &h            &     \\\\ \\hline \\hline 970402.930     &2.2$\\pm$0.6 &8   &1.6  &Ni98  &21    &18.5  &Gr97a \\\\ 971227.34938   &2.6$\\pm$0.6 &14  &1.12 &An99a  &22.8  &21.3  &Gr97b \\\\ 981226.40793   &5$\\pm$1     &11  &1.3  &Fr00  &23    &10.   &Li99  \\\\ 990217.22462   &$<1$     &6   &$>$1.6  &Pi99a &23.5  &19.   &Pa99a  \\\\ 990627.20894   &3.5         &8   &---  &Ni99b &21.   &23.   &Ro99  \\\\ 990704.7294    &4.4$\\pm$0.3 &8   &---  &Fe99  &22.5  &4.6   &Je99b  \\\\ 990806.60286   &5.5$\\pm$1.5 &7.8 &---  &Fr99  &22.   &3.8   &Vr99c \\\\ 990907.7319    &15$\\pm$5    &11  &---  &Pi99b &22.9$^a$ &24.9 &Pa99b \\\\ 990908.00125   &NP          &--- &---  &---   &20.$^b$ &11.5 &Ax99   \\\\ 991014.9115    &3.5$\\pm$0.5 &13 &$>$0.4 &Za00 &22.6  &12.9  &Ug99   \\\\ 991105.69495   &NP          &--- &---  &---   &23.5  &16.   &Pa99c   \\\\ 991106.4545   &1.25$\\pm$0.3 &8   &---  &An99b &21.   &9.1   &Ca99b   \\\\ 000210.36396   &4.5         &7.2 &---  &Co99b &23.3  &16.   &Go00a,b   \\\\ 000214.042    &2.75$\\pm$0.9 &12  &0.6  &An00  &21.$^c$ &32.4 &Rh00 \\\\ 000424.76258   &---         &--- &---  &---   &22.8  &33.   &Ug00   \\\\ 000528.36568   &1.7$\\pm$0.3 &8.3 &1    &Ku00  &23.3  &18    &Pa00a  \\\\ 000529.3361    &2.8$\\pm$0.7 &7.5 &---  &Fe00  &22.3  &47    &Pa00b   \\\\ 000615.2625    &---         &10  &---  &BS00  &21.5  &4.2   &St00   \\\\ 000620.2317    &---         &--- &---  &---   &19.8  &5.7   &Go00c   \\\\ \\hline 990520.08539   &---         &--- &---  &---   &21.7$^d$&19.5  &Ma99  \\\\ 991217.17496   &---         &--- &---  &---   &22.   &11.   &Mo99   \\\\ 000416.6062    &---         &--- &---  &---   &20.7  &50.3  &Pr00c   \\\\ \\hline \\hline \\end{tabular} \\\\ \\begin{flushleft} Notes:NP=Repointing of \\sax~not possible; $^a$: $R$ mag derived from V$>23.2$. $^b$: $R$ mag derived from $V>20.3$. $^c$: $R$ mag derived from $K>18.15$. $^d$: $R$ mag derived from $V>22$. \\\\ An99a: Antonelli et al., 1999a; An99b: Antonelli et al., 1999b; An00: Antonelli et al., 2000; Ax99: Axwlrod et al., 1999; BS00: \\sax~mail \\# 00/18 = GCN Circ. \\# 707; Ca99b: Castro-Tirado et al., 1999b; Co99b: Costa et al., 1999b; Fe99: Feroci et al., 1999; Fr99: Frontera et al., 1999; Fr00: Frontera et al., 2000; Go00a: Gorosabel et al., 2000a; Go00b: Gorosabel et al., 2000b; Go00c: Gorosabel et al., 2000c; Gr97a: Groot et al., 1997a; Gr97b: Groot et al., 1997b; Li99: Lindgren et al., 1999; Je99b: Jensen et al., 1999b; Ma99: Masetti et al., 1999; Mo99: Mohan et al., 1999; Ni98: Nicastro et al., 1998; Ni99b: Nicastro et al., 1998; Pa99a: Palazzi et al., 1999a; Pa99b: Palazzi et al., 1999b; Pa99c: Palazzi et al., 1999c; Pa00a: Palazzi et al., 2000a; Pa00b: Palazzi et al., 2000b; Pi99a: Piro et al., 1999a; Pi99b: Piro et al., 1999b; Pr00c: Price et al., 2000; Rh00: Rhoads eta l., 2000; Ro99: Rol et al., 1999; St00: Stanek et al., 2000; Ug99: Uglesich et al., 1999; Ug00: Uglesich et al., 2000; Vr99c: Vreeswijk et al., 1999c; Za00: in't Zand et al., 2000. \\end{flushleft} \\caption{{Properties of the bursts with \\sax-WFC detection but without associated optical transient. The last three bursts below the horizontal line refer to the $\\gamma$-ray poor GRBs (or X-ray transients) detected by \\sax.} \\label{tab:due}} \\end{table*} ",
        "conclusions": "By analyzing the properties of detected optical and X-ray afterglows and the upper limits for failed detections, we show that the subset of bursts without optical afterglow (FOAs) defines a different family. This conclusion relies on several assumptions, like the homogeneity of the \\sax~and non~\\sax~detected afterglows, imposed by the paucity of the sample. Should some of these turn out to be wrong, the conclusion would become statistically less stringent (see \\S2 for further details). We have investigated if this can be due to dust extinction of optical radiation in a molecular cloud.  We find that this hypothesis can only marginally account for the large fraction of FOAs, and therefore we cannot exclude the possibility that FOAs are intrinsically less luminous in the optical/UV band with respect to the detected ones, and with respect to their own X-ray luminosity.  Consider also that we have been very conservative in our procedure, because our results are based on considering {\\it upper limits} on the optical flux, and {\\it peak} absorption columns expected in giant molecular clouds.  The latter assumptions may well be too conservative, if the dust is bound to evaporate when illuminated and heated by the powerful optical/UV flash of the gamma-ray burst (Waxman \\& Draine 2000) and by its X-ray radiation (Fruchter et al. 2000).  This dust sublimation is suggested for a sample of burst afterglows (Vreeswijk et al. 1999c, Galama \\& Wijers 2000), in which a very large hydrogen column density $N_{\\rm H} \\gsim 10^{22}$ cm$^{-2}$, as estimated by X-ray data, is associated with almost no optical extinction.  The results can be understood only in terms of a dust to gas ratio $\\sim 100$ times smaller than the Galactic average value.  In turns, such low values of the dust to gas ratio can be explained only if the dust has been completely sublimated in the surroundings of the burst.  Indeed the theoretical models mentioned above predict that dust can be destroyed by the burst emission out to a radius comparable to the dimension of a typical molecular cloud (up to a few tens of parsecs).  If this is the case, the material responsible for absorption in FOAs is not the overdense cocoon surrounding the star forming region, but the cloud as a whole (or even less), and the discrepancy between the observed and measured value (see dash-dotted line in Fig.~\\ref{fig:mix}) becomes extremely compelling.   An interesting way to assess whether dust is playing any role in FOAs is to perform near infrared (NIR) follow-up of their $\\gamma$ or X-ray error boxes.  For instance, in the $K$ filter, absorption is greatly reduced, so that only a very small fraction (less than 10 per cent) of afterglows should show more than 1 magnitude of absorption, in any of the adopted cloud models.  This is therefore a crucial test to understand whether FOAs are due to dust absorption (less severe in the near infrared) or to an intrinsic difference in the emitted spectrum (that should be more severe in the NIR).  Some FOAs have been indeed looked for in the NIR band, but the observations are still very sparse and we lack any statistics to draw any meaningful conclusion.  NIR observations are thus strongly recommended as the key test for the dust extinction hypothesis, especially after the launch of HETE II, which will rapidly distribute accurate enough locations of bursts to be promptly followed by ground based telescopes.  A more homogeneous dataset, though, will have to await the launch of the Swift satellite, foreseen in 2003.  The systematic follow-up with the on-board optical telescope will provide a multiband spectroscopic database of the first hours of optical afterglows.  Data of even higher quality could be achieved if IR robotic telescopes (such as the one proposed by the consortium of Brera, Rome and Catania Observatories, called REM, for Rapid Eye Mount), will be in operation to complement Swift observations from the ground.  A possibility to increase the absorption in the observed $R$ band without invoking particularly dense molecular clouds is by allowing for a higher redshift of the bursts.  This would make the afterglow undetectable, especially if the redshift of the burst is particularly high ($z\\gsim 4$), so that the redshifted Lyman $\\alpha$ break falls in the $R$ filter.  In this case, again, near infrared ($JHK$) observations should be unaffected by absorption and the optical transient easily detectable. Estimates of the fraction of high redshift GRBs (see, e.g., Porciani \\& Madau 2001) predict however a very small fraction of bursts (up to few per cent) at $z>3$, if the GRB and star-formation rate are related. A more interesting way-out is that the property of clouds at high redshift are different from those of our Galaxy (see Ramirez-Ruiz Trentham \\& Blain 2001), or that the dust extinction curve changes its shape with redshift.  In conclusion, we have found that dust absorption due to a cloud with properties similar to Galactic clouds is not a completely satisfying explanation for bursts without a detected optical afterglow, and we cannot rule out the possibility that they are due to an intrinsic larger dispersion of optical fluxes with respect to the dispersion of the X-ray fluxes (see also B\\\"oer \\& Gendre 2000).  This, in turn, opens some exciting observational perspectives aiming to disclose the nature of the burst progenitor: if bursts are indeed associated with the final stages of stellar evolution and a supernova-like event is associated to all bursts, then the search for supernova signatures should be easier for bursts with an optical faint afterglow, for which the SN lightcurve would not be polluted by the flux of the afterglow.  If we assume SN1998bw (Galama et al. 1998c) as a template supernova lightcurve, the expected magnitudes at maximum should be roughly $I=24$ and $R=25$, easily detectable with a signal to noise ratio of $\\sim 10$ with an exposure time of only $\\sim 10$~min with an 8 meter class telescope."
    },
    "0011/astro-ph0011505_arXiv.txt": {
        "abstract": "HST V and I-band observations show that the gravitational lens B1359+154 consists of six images of a {\\it single} $z_s=3.235$ radio source and its star-forming host galaxy, produced by a compact group of galaxies at $z_l \\simeq 1$.  VLBA observations at 1.7 GHz strongly support this conclusion, showing six compact cores with similar low-frequency radio spectra. B1359+154 is the first example of galaxy-scale gravitational lensing in which more than four images are observed of the same background source. The configuration is due to the unique lensing mass distribution: three primary lens galaxies lying on the vertices of a triangle separated by $0\\farcs7 \\simeq 4 h^{-1}$ kpc, inside the $1\\farcs7$ diameter Einstein ring defined by the radio images. The gravitational potential has additional extrema within this triangle, creating a pair of central images that supplement the ``standard'' four-image geometry of the outer components. Simple mass models consisting of three lens galaxies constrained by HST and VLBA astrometry naturally reproduce the observed image positions but must be finely-tuned to fit the flux densities. ",
        "introduction": "The vast majority of the $\\sim 60$ known arcsecond-scale gravitational lens systems consist of two or four detectable images, consistent with the generic lensing properties of smooth, isolated and centrally steep mass distributions (e.g.\\ Blandford \\& Kochanek 1987). There are very few cases in which a ``non-standard'' number of images may have been observed.\\footnote{B1933+503 (Sykes et al.\\ 1998) and B1938+666 (King et al.\\ 1997) contain ten and six lensed radio components, respectively, but this is due to the imaging of a multi-component source rather than any exotic properties of the lensing potential.} APM08279+5255 (Ibata et al. 1999) and MG1131+0456 (Chen \\& Hewitt 1993) each contain a central component that could be an additional lensed image created by a sufficiently large galaxy core or shallow mass profile (Narasimha, Subramanian \\& Chitre 1986; Rusin \\& Ma 2000). Alternatively, APM08279+5255 may be a special class of imaging produced by an edge-on disk (Keeton \\& Kochanek 1998), while the central component in MG1131+0456 could be weak AGN emission associated with the lensing galaxy. In addition, MG2016+112 (Lawrence et al. 1984; Garrett et al. 1996) exhibits a rather complicated image morphology consisting of four primary components, one of which has three subcomponents in VLBI maps. This may be a compound lens produced by two galaxies at different redshifts (Nair \\& Garrett 1997).  Benitez et al.\\ (1999) suggest an alternative model, but it might not fully explain the VLBI observations.  Complex interplay between mass distributions can lead to lens systems with more than four images of a source. For example, ellipsoidal mass distributions perturbed by shear fields may produce configurations with six or eight images arranged about the tangential critical curve (Keeton, Mao \\& Witt 2000b). However, this requires that the relative magnitudes and orientations of the internal and external shear axes be finely tuned, and the resulting cross-sections are quite small. Compound mass distributions have been shown to be far more efficient at producing a variety of complex geometries in which five or more images may be formed over a range of radii (Kochanek \\& Apostolakis 1988).  This has been observed in the cluster-lensing case, where substructure in the gravitational potential created by the resident galaxies can qualitatively alter the lensing properties that one would expect for a smooth halo mass distribution (e.g. Natarajan et al. 1998; Meneghetti et al. 2000). A dramatic example of this is CL0024+1654, which exhibits eight images of a single blue background galaxy (Kassiola, Kovner \\& Fort 1992; Wallington, Kochanek \\& Koo 1995; Colley, Tyson \\& Turner 1996; Tyson, Kochanski \\& Dell'Antonio 1998).  Early-type galaxies preferentially participate in lensing (Kochanek et al.\\ 2000a), so a large fraction of lens galaxies should be members of small groups and clusters due to the morphology-density relation (Dressler 1980; see also Keeton, Christlein \\& Zabludoff 2000a). Indeed, many lens systems are known to be perturbed by nearby galaxies (e.g.\\ B2319+051; Rusin et al.\\ 2000b), groups (e.g.\\ PG1115+080; Schechter et al.\\ 1997), or clusters (e.g.\\ QSO 0957+561, Fischer et al.\\ 1997; RXJ0921+4529, Mu\\~noz et al.\\ 2000). A handful of systems are lensed by more than one primary galaxy (e.g. B1127+385, Koopmans et al.\\ 1999; B1608+656, Koopmans \\& Fassnacht 1999), but in these cases standard image geometries are produced despite merged caustics. Several additional lenses are observed to have faint satellite galaxies as companions inside or near the Einstein radius (e.g.\\ MG0414+0534, Schechter \\& Moore 1993; B1030+074, Xanthopoulos et al.\\ 1998; B1152+199, Rusin et al.\\ in preparation). While the likelihood of finding nearly equal mass galaxies close enough to have merged caustics is predicted to be only $\\sim 1\\%$ (Kochanek \\& Apostolakis 1988), the probability of finding fainter satellites near the primary lens is not small because galaxy luminosity functions diverge for faint galaxies -- all lenses should have faint neighbors, as discussed in the Appendix. While such systems typically produce regions of the source plane in which more than four images can form, the small size of the companion galaxies means that the cross-sections for creating non-standard image geometries are not significant.  The gravitational lens system B1359+154 (Myers et al.\\ 1999), discovered in the Cosmic Lens All-Sky Survey (CLASS; e.g.\\ Myers et al.\\ 1995), has been suspected of containing more than four lensed images. Observations with the Very Large Array (VLA) and Multi-Element Radio-Linked Interferometer Network (MERLIN) show a total of six radio components (Myers et al.\\ 1999; Rusin et al.\\ 2000a). Four of these components (A--D) are arranged in a typical quad-lens configuration (maximum image separation of $1\\farcs7$), with two additional components (E and F) residing within the ring defined by the outer images. Preliminary radio spectral studies of B1359+154 at high frequency (Myers et al.\\ 1999) suggested that E had a slightly flatter spectrum than A--D, and therefore that the central components may be core-jet emission associated with a weak AGN in the lensing galaxy or galaxies, as in the case of B2045+265 (Fassnacht et al.\\ 1999). When the spectra were extended down to 5 GHz, however, there appeared to be less disparity among the radio components (Rusin et al.\\ 2000a). Subsequent VLA observations at 15~GHz have failed to decisively confirm that E is flatter than A--D at high-frequency, or detect component F. Spectroscopy with the Keck II telescope determined the source redshift to be $z_s = 3.235$ (Myers et al.\\ 1999). Adaptive optics observations of B1359+154 conducted with the Canada-France Hawaii Telescope (CFHT) in the infrared K$'$ band (Rusin et al.\\ 2000a) detected counterparts to radio components A--D, and discovered three extended emission peaks (K1--K3) bracketing the expected positions of E and F. K1--K3 were identified as three possible lensing galaxies, comprising the core of a compact galaxy group. Evidence of an arc connecting A, B and C was also observed, along with a weaker emission feature associated with component E. The compound deflector system not only explained why attempts to model the outer four components using a single galaxy had failed (Myers et al.\\ 1999), but offered the means of creating additional extrema in the lensing potential. This opened the possibility that at least one of the central components is a lensed image.   In this paper we present powerful new evidence from observations with the Hubble Space Telescope (HST) and Very Long Baseline Array (VLBA) that B1359+154 consists of six images of a {\\em single} background source, lensed by a compact group of galaxies at $z_l\\simeq 1$. In \\S2 we present VLBA observations of B1359+154 and investigate the low-frequency radio spectra of the components. In \\S3 we discuss and analyze HST V and I-band observations, which offer compelling evidence for the six-image hypothesis. In \\S4 we use preliminary mass modeling to demonstrate that B1359+154 can be naturally explained as a true six-image lens system. Finally, in \\S5 we discuss the prospects for obtaining improved constraints on the lensing mass distribution of B1359+154, and ultimately using the system to study the structure of small galaxy groups at high redshift. ",
        "conclusions": "HST and VLBA observations demonstrate that the gravitational lens B1359+154 consists of six images of a single background radio source. This is the first example of a galaxy-scale lens system with more than four images.  The unique configuration is produced by the complex mass distribution of the lens, a compact group of three $z_l\\simeq 1$ galaxies lying on the vertices of a triangle separated by $0\\farcs7 \\simeq 4 h^{-1}$kpc, inside the $1\\farcs7$ diameter Einstein ring defined by the radio components. The outer images (A--D) have the morphology of a standard four-image lens, but the triangle of galaxies produces additional extrema in the gravitational potential at which two central images (E and F) form. Simple lens models consisting of three deflectors constrained by HST galaxy coordinates and VLBA radio component data naturally produce six images at the observed positions, but require finely-tuned ellipticities to account for the flux density ratios.  While it would be premature to claim that we have accurately represented the lensing mass distribution ($\\chi^2_{tot}$/NDF $\\geq 10$ for all models), our initial modeling efforts demonstrate that B1359+154 can be explained as a true six-image lens system produced by three lensing galaxies. The consistency between the model-predicted deflector mass ratios and the observed luminosity ratios of the lens galaxies adds credibility to our results.  The three primary lens galaxies appear to constitute the core of a compact group. The most concentrated groups with small numbers of galaxies are the Hickson groups, a quarter of which have 3 members (Hickson et al. 1992). Separated by only $\\simeq 4 h^{-1}$ kpc, the three galaxies are close enough so that they might be expected to be consumed in a merger on a time scale of $\\sim 0.1$ $H_o^{-1}$ (Barnes 1985). Thus the incidence of high multiplicity arcsecond-scale lens systems may allow for an estimate of the merger rate of massive galaxies at $z \\simeq 0.5 -1.5$.  The relative simplicity of B1359+154 compared to the lenses found in the cores of rich clusters should make this system an excellent tool for studying galaxy halos in dense environments, and the relationship between the galaxies and their parent halos. One issue worthy of investigation is whether the galaxy mass distributions are truncated on the scale of inter-galaxy separation. The improved fits of the 3PJS and 3PJE/FIX models relative to the 3SIS and 3SIE/FIX models offer tentative arguments to this effect. Because predicted time delays in the PJ models are significantly higher than those in the SIE models, measured delays may be used to discriminate between the profiles. A second interesting issue is the placement of the group halo relative to the lensing galaxies. The image positions and flux densities can be reproduced by three mass distributions associated with the observed galaxies, so there is little need to add a fourth independent mass distribution inside the Einstein ring based on modeling arguments. A fourth deflector above the critical density for multiple imaging would place an additional radial caustic onto the source plane and disturb the delicate balance needed to account for the data without creating additional images.  Thus if this halo exists within the Einstein ring, it is likely to be either subcritical (large core radius) or centered on one of the galaxies. However, the similar model-predicted mass-to-light ratios of the galaxies argue against the latter possibility. The large external shear fields ($\\gamma \\gtorder 0.15$) required by all of our models suggests that a group halo may be displaced from the primary lenses. It is interesting to note that the orientation of the shear axis is remarkably constant, running nearly east to west regardless of the deflector model employed.   The six lensed components allow for more complicated models than we have investigated here, but the development of a truly robust model of the lensing mass distribution requires the acquisition of additional constraints. The detection of lensed extended emission from the quasar host galaxy in both the HST optical and CFHT infrared images is particularly promising.  Deep, high-resolution HST imaging using either WFPC2 or NICMOS could investigate the properties of the Einstein ring in detail and provide vital new constraints on the potential structure (see Kochanek et al.\\ 2000b). Furthermore, if subcomponents of D, E and F can be detected using deep global VLBI imaging, the relative orientations of the radio substructure could place additional constraints on the local magnification matrices. Finally, the measurement of differential time delays would also offer important constraints on the mass model, as these directly probe the lensing potential at the image positions. One might expect the source to be variable, based on the compactness of the radio components, though there is no evidence for significant variability at this point. Radio snapshots of B1359+154 should be routinely obtained to determine whether the variability is large enough to undertake a monitoring program using current instruments. Otherwise, the enhanced VLA may be able to determine time delays from the component light curves even if the source variability is only at the few percent level.  It seems unlikely that the lens is a chance projection of three galaxies (an {\\it a posteriori} probability of $\\sim 10^{-3}$ in a sample of 60 lenses), and their statistically identical colors further argues against this possibility. While the Next Generation Space telescope could directly determine the redshifts of all three galaxies, current telescopes can only measure the average redshift of the lenses unless they have narrow emission lines that are not apparent in the Myers et al. (1999) spectrum. Two possibilities are to search for Mg{\\sc~II} metal absorption lines in the source spectrum or HI absorption features in the radio continuum once the average redshift is known.  For $z_l\\simeq1$ the Mg{\\sc~II} absorption features would lie in the V-band where the source flux is 23~V~mag, so the observation is difficult.  Many groups have significant HI masses (Hickson 1997) and the radio flux density of $\\simeq 90$~mJy at 0.7~GHz (based on the 1.7 GHz VLBA flux density and a spectral index of $\\alpha \\simeq 0.3$ at low frequency) may be high enough to search for absorption features.  X-ray observations are less promising because very deep observations would be required to detect the group (a bright $10^{42}$~erg~s$^{-1}$ group at $z_l=1$ has an X-ray flux of only $4\\times 10^{-16}$~ergs~cm$^{-2}$~s$^{-1}$) and the high resolution of the Chandra Observatory would be needed to distinguish emission from the group and the lensed sources.  \\noindent Acknowledgements: The authors deeply thank the director of STScI, Steve Beckwith, for making these observations possible.  We also thank Lars Hernquist and Ann Zabludoff for discussions about groups and the implications of this lens.  DR acknowledges support from the Zaccheus Daniel Foundation. Support for the CASTLES project was provided by NASA through grant numbers GO-8175, GO-8268 and GO-8804 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc.  CSK, EEF, JL, BAM and JAM were also supported by the Smithsonian Institution.  CSK and CRK were also supported by the NASA Astrophysics Theory Program grant NAG5-4062.  HWR is also supported by a Fellowship from the Alfred P. Sloan Foundation.   \\appendix"
    },
    "0011/hep-ph0011030_arXiv.txt": {
        "abstract": "If neutrino oscillation plays a role in explaining the atmospheric neutrino deficit, then the same phenomenon would necessarily affect also the dark matter indirect-detection signal which consists in a muon-neutrino flux produced by neutralino annihilation in the Earth core. In this paper we investigate to which extent the upgoing-muon signal originated by neutralinos captured inside the Earth would be affected by the presence of $\\nu_\\mu \\rightarrow \\nu_\\tau$ oscillation. ",
        "introduction": "Among the different techniques which have been proposed to search for dark matter particles \\cite{nota0}, detection of a neutrino flux by means of neutrino telescopes represents certainly an interesting tool. Although direct detection \\cite{ICTP,JKG,old} at present appears to be somewhat more sensitive to neutralino dark matter \\cite{probing,further,notadir}, nevertheless all the different possibilities are worth being explored.  In this paper we discuss the flux of upgoing muons which are a consequence of $\\nu_\\mu$'s produced inside the Earth, with special emphasis on the role which is played in this kind of searches by the possible presence of neutrino oscillation. ",
        "conclusions": "In this paper we have discussed to which extent neutrino oscillation can affect the up--going muon signal from neutralino annihilation in the Earth. While the experimental upper limit is, at present, practically not affected by neutrino oscillation \\cite{MACRO}, the theoretical predictions are reduced in the presence of oscillation. By adopting the neutrino oscillation parameters deduced {}from the fits on the atmospheric neutrino data \\cite{FGGV,SK_OSC_nos}, the effect is always larger for lighter neutralinos. For $\\nu_\\mu \\rightarrow \\nu_\\tau$ the reduction is between 0.5 and 0.8 for $m_\\chi \\lsim 100$ GeV and less than about 20\\% for $m_\\chi \\gsim 200$ GeV."
    },
    "0011/hep-ph0011176_arXiv.txt": {
        "abstract": "We review the phenomenology of UHE neutrino detection.  The  motivations for looking for such neutrinos, stemming from observational evidence and from the potential for new physics discoveries are enumerated, and their expected sources and   fluxes are given. Cross-sections with nucleons at energies all the way upto $10^{20}$ eV and the attenuation of fluxes in the Earth, both of which are physics issues important to their detection, are discussed. Finally, sample event-rates for extant and planned Water/Ice Cerenkov detectors are provided. ",
        "introduction": "It is now widely beleived that the recent Super Kamiokande (SK) result {\\cite{fuk,sobel}} of an anomaly in the flavour ratios and zenith angle dependance of the atmospheric neutrino flux provides the firmest signal yet of physics beyond the Standard Model (SM)\\footnote{The first hints of this anomaly were provided by atmospheric neutrino data from the IMB detector, \\cite{IMB}} . The significance of this can be gauged by the fact that a signal for such physics has been sought for in all the varied and intensive experimental probes that the many sectors  of the SM have been subjected to over more than  twenty-five years. That this signal comes from the neutrino sector of the theory  perhaps  assumes increased significance when considered in conjuntion with the fact that two other existing experimental anomalies which persist, the solar  deficit \\cite{solar} and the LSND result \\cite{lsnd} are also within the neutrino sector.  The detailed interpretation of these anomalies and their relation with each other  is still a matter of intense  and ongoing theoretical and experimental activity. However,  neutrino masses, mixings and consequent oscillations have provided the simplest framework for understanding the experimental results \\cite{smir}.  What may be said with a reasonable degree of conviction is that once their interpretation is clear, a hitherto unprecedented window into the nature of the physical theory beyond the SM will probably have been opened.  It is within this context, perhaps, that one should view the theoretical and experimental efforts that are ongoing in the area of ultra high energy (UHE) neutrino physics (for reviews see \\cite {gqrs,h1,r98}). If the physics excitement that we have uncovered in our study of low-energy neutrinos is an indicator, then the study of UHE neutinos should pay rich dividends.  In view of the fact that data collection and further upgradation is ongoing for the first generation of UHE neutrino detectors ( AMANDA \\cite{aman} and  BAIKAL \\cite{bai}) and that planning, design, testing and deployment is underway for some others (AUGER \\cite {aug}, NESTOR \\cite{nest}, ICECUBE \\cite{ice},  ANTARES \\cite{ant},   RICE \\cite{rice}), and NEMO \\cite{nemo}) definitive results of these efforts are still in the future. However, besides the other  specific reasons from astrophysics and cosmic ray physics (some of which we discuss below) which provide motivation for such experiments, we note that the energy range covered by these experiments ($E_{\\nu} \\simeq 1$ Tev to $10^9$ TeV or higher) offers an unprecedented opportunity for particle physics at energies significantly beyond the scope of terrestial accelerators. Although progress is expected to be slow, the potential for serendipitous discovery is undoubtedly high.  In what follows, we  give specific reasons why the search for UHE neutrinos can be expected to yeild positive results. Possible sources and their fluxes are discussed. Salient features pertinent to the phenomenology of  detection are then outlined, and sample rates provided. ",
        "conclusions": "We have made an attempt to review the  essential motivations and  phenomenological issues related to the interactions and detection of UHE neutrinos and provided expected event rates for water/ice Cerenkov detectors.  Although present detector capabilities and sizes do not allow them to see above atmospheric backgrounds, with upgradations and  several new  large-scale experiments underway, the detection of UHE neutrinos from astrophysical sources may be very likely in the near future. Besides the potential for serendipitous discovery, the unmatched (terrestially) energies may lead to new particle physics discoveries. In addition, these experiments may help clarify the astrophysical mechanism responsible for the ultra-relativistic acceleration of matter in these sources and answer important questions in CR physics, and, finally, provide sensitive tests of gravitational couplings and oscillations. For those working in UHE neutrino physics and astronomy, the anticipation is palpable."
    },
    "0011/astro-ph0011263_arXiv.txt": {
        "abstract": "Here we report the detection, in H$\\alpha$ emission, of a radial corrugation in the velocity field of the spiral galaxy NGC 5427. The central velocity of the H$\\alpha$ line displays coherent, wavy-like variations in the vicinity of the spiral arms. The spectra along three different arm segments show that the maximum amplitude of the sinusoidal line variations are displaced some 500 pc from the central part of the spiral arms. The peak blueshifted velocities appear some 500 pc upstream the arm, whereas the peak redshifted velocities are located some 500 pc downstream the arm. This kinematical behavior is similar to the one expected in a galactic bore generated by the interaction of a spiral density wave with a thick gaseous disk, as recently modeled by Martos \\& Cox (1998). ",
        "introduction": "It has been known for decades that the gaseous disk of our Galaxy displays coherent vertical distortions (Kerr 1957; Gum, Kerr \\& Westerhout 1960). In particular, for locations inside the solar circle, a wavy structure has been observed in the vertical distribution of several interstellar tracers, including the atomic and ionized components (Lockman 1977; Sanders, Solomon \\& Scoville 1984; Spicker \\& Feitzinger 1986; Malhotra 1995), and young stellar objects (\\eg Dixon 1967; Alfaro \\etal 1992a,b; Berdnikov \\& Efremov 1993). These undulations, distributed above and below the mean Galactic plane and in both the azimuthal and radial directions, have been termed as ``corrugations'', and they seem to be common features appearing in disk galaxies. They have also been detected in both the gaseous and stellar disks of other spiral galaxies such as M31 (Arp  1964), NGC 4244, and NGC 5023 (Florido \\etal 1991).  Their origin could be ascribed to either well localized features, such as spiral arms and collisions with high-velocity clouds (HVC), or large scale perturbations, such as gravitational interactions. For instance, some authors have explored the possibility of undulations along spiral arms, induced by magneto-gravitational instabilities (\\eg Nelson 1985; G\\'omez de Castro \\& Pudritz 1992; Kim \\etal 1997; Franco \\etal 2000). In this case, the arms are corrugated by the undular mode (Parker 1966) of the instability, and the structuring occurs $only$ in the azimuthal direction (a possible association to the radial patterns could be made if one assumes that these are the projection of the azimuthal corrugations on the radial direction). Another localized origin could be ascribed to collisions of HVC with the disk of a galaxy: HVC impacts could cause the midplane disk to oscillate in a mode of wavy patterns as well (Franco \\etal 1988; Santill\\'an \\etal 1999). In the large scale scenario, the undulations represent the natural response of the disk, which behaves like a loose drumhead, to the perturbations induced either by some companion galaxies or other galactic subsystems (Weinberg 1991; Edelsohn \\& Elmegreen 1997).  One obvious inference is that these corrugations should be associated with equally undulated coherent motions, and models addressing the vertical displacements also show wave-like $z$-velocity fields with amplitudes of the order of tens of km~s$^{-1}$ (\\eg Nelson 1976, 1985; Franco \\etal 1989, 2000; Edelsohn \\& Elmegreen 1997; Kim \\etal 1997; Santill\\'an \\etal 2000). Despite this expectation, however, the data about actual velocity features associated with spatial corrugations are scarce. The observational evidence for small-scale distortions in the vertical velocity fields of disk galaxies comes from the so-called {\\sl rolling motions} in our Galaxy (observed in the $b - V$ plots of the HI galactic distribution; Yuan \\& Wallace 1973; Feitzinger \\& Spicker 1985), the velocity field of the ionized gas in M 51 (Goad, De Veny \\& Goad 1979), rotation curves of some spirals (Pismis 1965; Rubin, Ford \\& Thonnard 1980), and optical and radio isovelocity maps of some field galaxies (Bosma 1978; Pismis 1986). In particular, the study of the velocity field of the ionized gas in M 51 performed by Goad, De Veny \\& Goad (1979) suggests that part of the observed nonrotational velocities could be due to motions perpendicular to the plane of the disk. Pismis (1965, 1986), on the other hand, has pointed out the relevance of analyzing the location of the velocity undulations in the rotation curve with respect to the spiral structure of the host galaxy. Clearly, all this kinematic information may be important in our understanding of corrugations.  In this paper we report the detection of a corrugated velocity field in the ionized gas of the nearly face-on spiral NGC 5427. The central velocity of the H$\\alpha$ emission line shows velocity shifts with similar behavior when the gas passes through the spiral arms. This behavior can be understood as due to a hydraulic jump, or bore, generated by the interaction of the spiral density wave with the gas of a thick gaseous disk, as discussed by Martos \\& Cox (1998, hereafter MC) and Martos \\etal (1999). The detection of this velocity pattern in NGC 5427 does not answer the current questions about the morphogenesis of corrugations, but represents a powerful tool to explore further the theoretical and observational aspects of the problem. The paper has been organized in three sections, the first one being this introduction; section 2 is devoted to the description of the observations and reduction procedure and, finally, in section 3, we discuss the results and outline the main conclusions. ",
        "conclusions": ""
    },
    "0011/astro-ph0011055_arXiv.txt": {
        "abstract": "\\noindent The observational data on the large scale structure (LSS) of the Universe are used to establish the upper limit on the neutrino content marginalized over all other cosmological parameters within the class of adiabatic inflationary models. It is shown that the upper 2$\\sigma$ limit on the neutrino content can be expressed in the form $\\Omega_{\\nu}h^2/N_{\\nu}^{0.64}\\le0.042$ or, via the neutrino mass, $m_{\\nu}\\le4.0$eV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011325_arXiv.txt": {
        "abstract": "We present a new analysis of the currently available orbital elements for the known Kuiper belt objects.  In the non-resonant, main Kuiper belt we find a statistically significant relationship between an object's absolute magnitude ($H$) and its inclination ($i$).  Objects with $H<6.5$ (i$.$e$.$ radii $\\gapprox 170$km for a 4\\% albedo) have higher inclinations than those with $H>6.5$ (radii $\\lapprox 170\\,$km).  We have shown that this relationship is not caused by any obvious observational bias.  We argue that the main Kuiper belt consists of the superposition of two distinct distributions.  One is dynamically hot with inclinations as large as $\\sim 35^\\circ$ and absolute magnitudes as bright as $4.5$; the other is dynamically cold with $i\\lapprox 5^\\circ$ and $H>6.5$.  The dynamically cold population is most likely dynamically primordial.  We speculate on the potential causes of this relationship. ",
        "introduction": "\\label{sec_intro}  The discovery of the Kuiper belt in 1992 (Jewitt \\& Luu~1993) issued in a new era for the study of the outer solar system.  The Kuiper belt is important not only because it is a rich, new region of the solar system to be explored, but because it contains important fossil clues about the formation of the outer solar system in particular, and about planet formation in general.  Since its discovery, the Kuiper belt has supplied us with surprise after surprise.  For example, before it was discovered, theorists believed that the Kuiper belt would consist of objects on low-inclination, nearly-circular orbits beyond the orbit of Neptune (Levison \\& Duncan~1993; Holman \\& Wisdom~1993).  This belief seemed to be confirmed with the discovery of the first two Kuiper Belt Objects (hereafter KBOs), 1992~QB$_1$ and 1993~FW.  However, the next four objects discovered revealed a real surprise.  At the time of discovery their heliocentric distances were close enough to Neptune's orbit that their orbits should be unstable, unless protected by some dynamical mechanism.  Indeed, many believed that they might have been Neptunian Trojans.  However, these were the first discoveries of an unexpected population of objects on highly eccentric (up to 0.3) orbits in the 2:3 mean motion resonance with Neptune (co-orbiting with Pluto).  Currently, objects in the trans-Neptunian region are divided into two main groups (see Malhotra et al$.$~2000 for a review).  The {\\it Kuiper belt} consists of objects that are primarily on long-lived orbits, while the {\\it scattered disk} consists of objects that have suffered a close encounter with Neptune (Duncan \\& Levison~1997; Luu et al$.$~1997).  The Kuiper belt itself is typically subdivided into two populations.  Inside of roughly $42\\au$, objects tend to be locked into mean motion resonances with Neptune.  Most known objects in this class are in Neptune's 2:3 mean motion resonance.  However, a fraction also reside in the 3:5 and the 3:4 resonances.  The orbits of all these objects are probably a result of resonance capture during the slow outward migration of Neptune during the late stages of planet formation (Malhotra~1995).  Beyond $42\\au$, although several objects are believed to be in the 1:2 mean motion resonance (Marsden~2000a), most objects are not on resonant orbits.  These non-resonant objects are members of what has come to be called the {\\it main Kuiper belt}.  Models of planetary migration (e$.$g$.$ Malhotra~1995; Holman~1995; Hahn \\& Malhotra~1999) predict that unlike the KBOs in mean motion resonances, main KBOs should be on relatively low-inclination, nearly-circular orbits. However, recent observations have shown that this is not the case. Numerous objects in this region have very large inclinations\\footnote{Eccentricities are not a good measure of how excited the Kuiper belt is since most large eccentricity orbits are removed through close encounters with Neptune, truncating the eccentricity distribution.  Inclinations do not suffer from this problem (Duncan, et al$.$~1995).}, certainly up to about $32^\\circ$, and most likely even higher (Marsden~2000a).  Several papers have been published which attempt, among other things, to explain the high inclinations seen in the main Kuiper belt. The mechanisms invoked to date involve the scattering of KBOs by large objects temporarily evolving through the region.  It takes a massive object to excite KBOs to high inclination; much more massive than the KBOs themselves\\footnote{A simple calculation based on an object's escape velocity shows that it must be larger than roughly twice the radius of Pluto to scatter a Kuiper belt object to an inclination of 30$^\\circ$.}.  Petit et al$.$~(1999) suggest that the dynamically excited Kuiper belt is caused by the passage of Earth-mass objects through that region of the solar system.  Thommes et al$.$~(1999) suggest that the large inclinations are due to the passage of Uranus and/or Neptune through the Kuiper belt while on eccentric orbits, after these planets were ejected from the region between Jupiter and Saturn.  Ida et al$.$~(2000) suggest that the Kuiper belt was excited by a passing star.  In this paper we present an analysis of the currently available orbital data of main belt KBOs which shows a new and surprising trend --- an unexpected and intriguing correlation between inclination and absolute magnitude.  In particular intrinsically bright objects tend to be found on larger inclinations than do intrinsically faint objects.  In \\S{\\ref{sec_data}} we present the data and discuss the statistical significance of this trend.  In \\S{\\ref{sec_bias}} we investigate whether this trend is a result of observational selection effects. Our preliminary interpretation of this trend is presented in \\S{\\ref{sec_interp}}. We summarize our findings in \\S{\\ref{sec_concl}}.  \\clearpage ",
        "conclusions": "\\label{sec_concl}  We have shown that the inclination distribution of objects in the main Kuiper belt most likely varies as a function of absolute magnitude. In particular, objects intrinsically brighter than $H=6.5$ appear to have systematically higher inclinations than intrinsically fainter objects.  There is only $\\sim 3\\%$ chance that these two distributions are the same.  We have shown that this result is unlikely to be caused by biases in discovery or recovery observing procedures.  Therefore, although it is possible that this conclusion is a result of small number statistics, we believe that it is real.  Future discoveries and followups will clearly resolve this issue.  The clear implication of our result is that a main belt object's inclination is dependent on its size.  The differences between intrinsically bright objects and the intrinsically faint objects is best seen in Figure~\\ref{fig_iH}. Perhaps the most natural interpretation for the data in this figure is that we are seeing the superposition of two distinct populations.  The first contains a dynamically hot population (inclinations up to $\\sim 35^\\circ$) consisting of both large and small objects (absolute magnitudes as small as 4.5 or radii up to $\\sim 330\\,$km for albedos of 4\\%).  Indeed, even larger objects and/or objects with higher inclinations are likely to still be found.  The other population is a dynamically cold one ($i\\lapprox 5^\\circ$) preferentially containing smaller objects ($H\\gapprox6.5$ or radii $\\lapprox 170\\,$km for albedos of 4\\%)."
    },
    "0011/astro-ph0011113_arXiv.txt": {
        "abstract": "We develop a method for recovering the global density distribution of the ancient Galactic stellar halo prior to disk formation, based on the {\\it present} orbits of metal-poor stars observed in the solar neighborhood. The method relies on the adiabatic invariance of the action integrals of motion for the halo population during the slow accumulation of a disk component, subsequent to earlier halo formation. The method is then applied to a sample of local stars with [Fe/H]$\\le-1.5$, likely to be dominated by the halo component, taken from Beers et al.'s recently revised and supplemented catalog of metal-poor stars selected without kinematic bias.  We find that even if the Galactic potential is made spherical by removing the disk component in an adiabatic manner, the halo density distribution in the inner halo region ($R \\le 15$ kpc) remains moderately flattened, with axial ratio of about 0.8 for stars in the abundance range [Fe/H]$\\le -1.8$ and about 0.7 for the more metal-rich interval $-1.8<$[Fe/H]$\\le-1.5$.  The outer halo remains spherical for both abundance intervals.  We also find that this initial flattening of the inner halo is caused by the anisotropic velocity dispersions of the halo stars. These results suggest that the two-component nature of the present-day stellar halo, characterized by a highly flattened inner halo and nearly spherical outer halo, is a consequence of {\\it both} an initially two-component density distribution of the halo (perhaps a signature of dissipative halo formation) {\\it and} of the adiabatic flattening of the inner part by later disk formation.  Further implications of our results for the formation of the Galaxy are also discussed, in particular in the context of the hierarchical clustering scenario of galaxy formation. ",
        "introduction": "The currently favored cold dark matter (CDM) theory of galaxy formation postulates that the formation of a massive spiral galaxy like our own is a consequence of the hierarchical assembly of subgalactic dark halos, and the subsequent accretion of cooled baryonic gas in a virialized, galaxy-scale dark halo (e.g., Peacock 1999). Numerical studies based on this picture are able to, at least qualitatively, reproduce the characteristic features of a disk galaxy -- the massive dark halo, the stellar halo, and the stellar disk components (e.g., Steinmetz \\& M\\\"uller 1995; Bekki \\& Chiba 2000; Navarro \\& Steinmetz 2000), though difficulties are still encountered in the details.  For example, the simulations conducted to date do not adequately account for the size of the disk component and the number of satellite galaxies (Navarro, Frenk, \\& White 1995; Moore et al. 1999; Klypin et al. 1999).  The CDM hierarchical model may be regarded, in its essence, as a modern generalization of the classical Searle \\& Zinn (1978) hypothesis for the formation of the Galactic stellar halo.  To explain a large inferred spread in the ages of globular clusters, and the lack of a spatial gradient in their metal abundances, Searle \\& Zinn argued that the halo component may have experienced prolonged, chaotic accretion of subgalactic fragments, as opposed to the rapid, monolithic collapse proposed by Eggen, Lynden-Bell, \\& Sandage (1962). Recent discoveries of halo substructures in Galactic phase space (Majewski, Munn, \\& Hawley 1994; 1996; Helmi et al. 1999; Chiba \\& Beers 2000, hereafter CB; Yanny et al. 2000) and of the Sagittarius dwarf galaxy, which is presently being disrupted by the Galactic tidal field (Ibata, Gilmore, \\& Irwin 1994; Ibata et al. 2000), may lend further support to this picture.  Although halo formation via hierarchical assembly of subgalactic systems, such as dwarf galaxies, may continue to the present day (Bland-Hawthorn \\& Freeman 2000), a large fraction of the stellar halo, especially the inner part where the disk lies, should have been completed {\\it prior to disk formation}, since otherwise the disk component is made significantly thicker than is observed due to dynamical heating from infalling masses (Toth \\& Ostriker 1992). A clear age gap between the (thin) disk and the stellar halo supports that the latter consists of ancient populations (e.g. Liu \\& Chaboyer 2000). Also, studies of star-forming histories in the disk component indicate that the disk has been accumulated at an approximately constant rate over the last several billion years (Twarog 1980; Sommer-Larsen \\& Yoshii 1990); frequent mergings of dwarf galaxies over the Galaxy's lifetime will entirely modify the photometric and spectroscopic properties of the disk. Thus, one may well postulate that the inner part of the stellar halo, at say $R \\le 15$ kpc, retains a fossil imprint of how it was formed.  The question then arises, what {\\it was} the structure of the halo component prior to disk formation?  Since the bulk of halo stars are found in this inner part, where the disk gravity is dominant, the present-day structure of the halo can be greatly affected by later disk formation.  It is thus necessary to consider the dynamical effect of the disk in inferring the structure of the ancient halo from the currently observed halo stars.  Binney \\& May (1986, hereafter BM) examined this issue by assuming adiabatic invariance (for halo stars) of the action integrals of motion ${\\bf J}$, and for the distribution function $f({\\bf J})$ during slow disk formation. They set up test particles distributed in a spheroid, similarly to the stellar distributions of elliptical galaxies, and calculated the dynamical response of the particles to the slow increase of disk mass inside the spheroid.  They showed that the Galactic halo, before the disk was formed, may have had a somewhat flattened density distribution (axial ratio $q \\sim 0.7$), in order to produce a current highly flattened halo ($q \\sim 0.3$), which they inferred from the radially anisotropic velocity ellipsoid of local metal-poor stars.  We note here that recent kinematic data for larger samples of local metal-poor stars indicate a more moderately flattened halo ($q \\sim 0.7$) for inner radii ($R \\le 15$ kpc), whereas the outer halo is nearly spherical (Sommer-Larsen \\& Zhen 1990, hereafter SLZ; CB). This is also supported by examination of the spatial distributions of other halo tracers (e.g., Hartwick 1987; Preston et al.  1991; Kinman, Suntzeff, \\& Kraft 1994; Yanny et al. 2000).  Also, the extent to which the formation of the disk component flattened the halo depends on the unknown initial velocity distribution of halo stars, whereas the initial conditions set up by BM apply only for one specific case.  Thus, it is yet unexplored what the currently available data for metal-poor stars may tell us about the structure of the ancient halo before the disk was formed.  In this paper we revisit this issue, based on a large sample of halo stars in the solar neighborhood, taken from a recently completed catalog of metal-poor stars selected without kinematic bias (Beers et al. 2000). It is noted that in a similar vein, Sommer-Larsen (1986), in his thesis work, arrived at a conclusion similar to BM's by investigating the distribution of individual orbital inclinations for his sample of 143 stars with [Fe/H]$\\le-1.2$. In contrast, we seek herein to develop a more general and direct method, based on the BM picture, to calculate the global density distribution of the halo prior to disk formation. The method is then applied to the more accurate and numerous data for metal-poor stars that is presently available.  This paper is organized as follows. In \\S 2 we describe general properties of orbits in the mass model of the St\\\"ackel type that we adopt here, as well as the methodology for constructing the global density of a given local sample both before and after disk formation. The mass model for the Galactic potential that we adopt, and the local sample of metal-poor stars used in the current analysis, are also described.  In \\S 3 we compute the orbital motions of the sample stars while conserving action integrals of motion.  We then present the adiabatic change of the derived density distributions and kinematic properties of the sample when the disk component is slowly removed (essentially working backwards from the present to the past).  Finally, in \\S 5, the results are summarized, and their implications for the formation and evolution of the Galaxy are discussed. ",
        "conclusions": "The global structure of the present-day stellar halo is characterized by an inner, highly flattened part, as revealed at $R<15$ kpc, and an outer, nearly spherical part (SLZ; CB). This two-component picture for the present-day stellar halo provides a reasonable explanation why faint-star-count studies have generally yielded an approximately spherical halo, whereas the local anisotropic velocities of the halo stars suggest a highly flattened system (Freeman 1987).  The issue relevant here is what physical mechanism in the early stage of the Galaxy gives rise to the inner, highly flattened halo, where the bulk of halo stars are found, and where the effects of later satellite accretion may be diminished.  One of the possible reasons for the two-component nature of the present-day halo is the slow formation of the disk within an initially spherical stellar halo (BM). In this paper, we have quantified the effect of later disk formation on the halo flattening, based on methods assuming adiabatic invariance of the motion of halo stars, and its application to a large sample of stars in the solar neighborhood. We have found that, even before disk formation, the inner part of the stellar halo exhibited a finite flattening, although it is more moderate than presently observed. The axial ratios of the density profiles within the almost spherical potential are $q \\simeq 0.80$ for [Fe/H]$\\le-1.8$ and $q \\simeq 0.70$ for $-1.8<$[Fe/H]$\\le-1.5$. Also, the initial velocity dispersions are characterized by an anisotropic velocity ellipsoid, as $(\\sigma_\\phi/\\sigma_\\lambda,\\sigma_\\nu/\\sigma_\\lambda)=(0.77,0.60)$ for [Fe/H]$\\le-1.8$ and (0.59,0.44) for $-1.8<$[Fe/H]$\\le-1.5$ at $R=R_\\odot$. Therefore, the inner part of the stellar halo was flattened by velocity anisotropy, and the more metal-rich population likely exhibited a more flattened density distribution.  Through a comparison of the inclination angles of orbits for two abundance ranges, [Fe/H]$\\le-1.5$ and $-1.5<$[Fe/H]$\\le-1.2$, Sommer-Larsen (1986) also obtained a more flattened initial distribution for more metal-rich populations. We note that his latter subsample shows a rapid systematic rotation of 147 km s$^{-1}$, possibly contaminated by the stars belonging to the metal-weak thick disk (Freeman 1987). In contrast, our metal-rich halo subsample is selected from the more restrictive range $-1.8<$[Fe/H]$\\le-1.5$, where the effect of disk-like kinematics is minimal (Chiba \\& Yoshii 1998; CB) -- the flattening of this subsample is caused by velocity anisotropy, not systematic rotation.  The results presented here may suggest that the ancient halo, at least in its inner part, may have undergone a somewhat ordered contraction.  This contraction could have involved dissipation due to baryonic gas -- radiative cooling of this gas was most efficient in the innermost regions with high density.  The resultant contraction of this gas, as a whole, may have proceeded mainly along the axis of rotation, because of the absence of the angular-momentum barrier in this direction.  As the chemical enrichment proceeded along with the progress of the collapse, more metal-enriched stars would have ``seen'' a more flattened density distribution.  On the other hand, the outer part of the halo may have been more susceptible to later infall of satellite galaxies, so that its density, kinematics, and mean age are different from those of the inner halo (Norris 1994; Carney et al. 1996; Sommer-Larsen et al. 1997).  Alternatively, one might argue that the two-component nature of the present-day stellar halo is entirely a consequence of satellite accretion {\\it after} disk formation (Freeman 1987). According to Quinn \\& Goodman (1986) (see also Quinn, Hernquist, \\& Fullagar 1993), prograde satellite orbits that are initially inclined at less than about 60$^\\circ$ to the disk are dragged down quickly toward the plane by the effects of the dynamical friction of the disk. Walker, Mihos, \\& Hernquist (1996) further explore the effects of mergers of small satellites with large disk galaxies such as the Milky Way.  The debris from these merging satellites, in combination with disrupted disk stars, would be expected to form a flattened system.  Furthermore, if more massive satellites were more metal-rich (as they appear to be at present, see Mateo 1998), their orbits would fall farther toward the Galactic center so that their debris would form a more flattened, more metal-rich system (Freeman 1987). However, there exists no clear evidence for the predicted dynamical heating of the thin-disk component -- its very thin geometry (Toth \\& Ostriker 1992), and the nearly constant velocity dispersion of thin-disk stars over the last 10 Gyrs (Quillen \\& Garnett 2000), suggest that the thin disk has sustained little significant damage since its formation.  In this regard, one might argue that the metal-weak thick disk is evidence for early dynamical heating of a pre-existing, metal-deficient, thin disk. However, it is then difficult to explain its absence in the abundance range considered in this paper (Chiba \\& Yoshii 1998; CB).  It is worthwhile to remark that the hypothesis of the dissipative formation of the inner flattened halo, as well as the later accretion of satellites onto the outer halo, is a natural consequence of the CDM hierarchical clustering model (Bekki \\& Chiba 2000).  This model postulates that a protogalactic system initially contains numerous subgalactic clumps, comprised of a mixture of gas and dark matter, and that the merging of these clumps led to a smaller number of more massive clumps.  In the simulations of Bekki \\& Chiba (2000), these larger clumps move gradually toward the central region of the system, due to both dynamical friction and dissipative merging with smaller clumps.  Finally, the last merging event occurs between the two most massive clumps, and the metal-poor stars which have been formed inside the clumps are disrupted and spread over the inner part of the halo.  The aftermath is characterized by a flattened density distribution.  Some fraction of the disrupted gas from the clumps may settle into the central region of the system, and produce a more enriched, more flattened density distribution.  Some of the initially small density fluctuations in the outer region would have gained systematically higher angular momentum from their surroundings, and then slowly fallen into the system after most parts of the system were formed.   This may correspond to the process of late satellite accretion, contributing primarily to the outer part of the halo.  Thus, the reported initial state of the stellar halo can be explained, at least qualitatively, in the context of hierarchical clustering scenario.  An alternative approach for elucidation of the dissipative nature of halo formation is to examine the results of recent high-resolution N-body simulations of structure formation based on the CDM theory (e.g., Ghigna et al. 1998; Moore et al. 1999; Klypin et al. 1999). Such simulations provide the orbital properties of dark matter particles inside virialized dark halos.  If the stellar halo component in the Galaxy is formed similarly through dissipationless hierarchical assembly, the orbits of halo stars prior to disk formation, as derived in the current paper, may follow those of dark matter particles. For this purpose, we take the Ghigna et al. (1998) simulation of the formation of a cluster, as this is currently the only published one that presents the detailed orbital distribution of the simulated particles.  They showed that the orbital distribution of the halo particles is close to isotropic -- circular orbits are rare and radial orbits are common.  The average ratio of pericentric and apocentric distances, $r_{pr}/r_{ap}$, is equal to or less than 0.20, without showing a large variation as a function of the distance from the cluster: the median ratio is approximately 0.17. On the other hand, as is deduced from Figure 2, the orbits of the halo stars we have derived here before the disk was in place are more circular than the simulated dark halo particles, and the velocity field is anisotropic: the average value of $r_{pr}/r_{ap}$ is 0.29.  Thus, we require some additional process, possibly dissipative interaction among protogalactic clumps to circularize their orbits, to explain the characteristic orbital distribution of the halo stars in the early Galaxy.  More quantitative conclusions must await more elaborate modeling of the formation of the Galaxy over a large number of possible model parameters. Also, it is necessary to assemble and analyze the data of more remote stars, especially those presently found inside the solar radius, where our modeling of the halo is incomplete. In this regard, the next generation of astrometric satellites, such as {\\it FAME} and {\\it GAIA}, will provide highly precise parallaxes and proper motions for numerous stars, so that both three dimensional positions and velocities will be available over a large fraction of the halo.  Also, with these astrometric satellites, we will be able to determine the exact mass distributions of the disk and dark halo components, and thus obtain definite information on the early Galaxy, using the technique outlined here.  Furthermore, in addition to the Milky Way, direct identification of halo populations in external disk galaxies may prove promising as a way to clarify the global structures of stellar halos and their association with disks and bulges (e.g., Morrison 1999). Such studies should be eagerly pursued with 10m class telescopes."
    },
    "0011/astro-ph0011439_arXiv.txt": {
        "abstract": "I consider the growth of inhomogeneities in a low-density baryonic, vacuum energy-dominated universe in the context of modified Newtonian dynamics (MOND). I first write down a two-field Langrangian-based theory of MOND (non-relativistic), which embodies several assumptions such as constancy of the MOND acceleration parameter, association of a MOND force with peculiar accelerations only, and the deceleration of the Hubble flow as a background field which influences the dynamics of a finite size region.  In the context of this theory, the equation for the evolution of spherically symmetric over-densities is non-linear and implies very rapid growth even in a low-density background, particularly at the epoch when the putative cosmological constant begins to dominate the Hubble expansion.  Small comoving scales enter the MOND regime earlier than larger scales and therefore evolve to large over-densities sooner.  Taking the initial COBE-normalized power spectrum provided by CMBFAST (Seljak \\& Zeldarriaga 1996), I find that the final power-spectrum resembles that of the standard $\\Lambda$CDM universe and thus retains the empirical successes of that model. ",
        "introduction": "A primary motivation for cosmic non-baryonic dark matter with negligible pressure is the necessity of forming the presently observed structure in the Universe without violating the constraints on temperature fluctuations in the CMB.  Basically, this is because structure in the dark matter component on galaxy to super-cluster scales can begin growing via gravitational instability considerably before hydrogen recombination (Peebles 1982, Vittorio \\& Silk 1984, Bond \\& Efstathiou 1984).  This remains one of the powerful arguments against a low density baryonic universe.  Any alternative cosmology not including CDM must invoke some mechanism other than conventional gravitational collapse in order to form structure.  McGaugh (1999) suggests that the modified Newtonian dynamics (MOND), proposed by Milgrom (1983) as an alternative to dark matter on on galaxy and cluster scales, can provide the needed mechanism and, further, that the consistency of the observed angular structure of the temperature fluctuations in the CMB (Lange et al. 2000, Hanany et al. 2000), with a pure baryonic universe (McGaugh 2000), may be viewed as support for MOND. This speculation is based upon the general expectation that MOND, in providing stronger effective gravity in the limit of low accelerations, would assist in structure formation.  MOND is an {\\it ad hoc} modification of Newton's law of inertia or gravity at low acceleration.  The original idea is contained in the statement that when the acceleration falls below $a_o$, a new physical constant with units of acceleration, then the effective gravitational acceleration approaches $\\sqrt{g_na_o}$ where $g_n$ is the usual Newtonian gravitational acceleration.  Although this simple formula works remarkably well in describing galaxy rotation curves consistently with the observed distribution of detectable matter (Sanders 1996, McGaugh and de Blok 1998), it clearly lacks the generality to treat the problem of cosmological density fluctuations.  A more consistent physical description of modified dynamics is provided by the non-relativistic Langrangian-based theory of Bekenstein and Milgrom (1984, hereafter BM).   An obvious procedure, when treating the growth of density fluctuations, would be to take the modified Poisson equation of BM and consider small fluctuations about a zeroth order solution as in Newtonian cosmology.  The problem is that, when applied to a finite sphere as is usual in Newtonian cosmology, the zeroth order solution is not that of a linear Hubble flow-- the absolute distance cannot be factored out and it is not possible to describe cosmology in terms of a universal scale factor.  The cosmology is basically that described by Felten (1984) and Sanders (1998), in which MOND alters the usual Friedmann solutions; as soon as the cosmic deceleration over some physical scale falls below $a_o$, then that  entire region begins to deviate from uniform Hubble flow.  This leads to the eventual re-collapse of any finite size region regardless of its original density or expansion velocity.  In this picture density fluctuations play no role.  Apart from problems in principle (what determines the point or points about which MOND collapse proceeds?), this cosmology leads to clear contradictions with observations-- recollapse in the present Universe occurs out to scales of 30 Mpc.  One might expect that in a  proper theory, the basic Hubble flow remains intact, and structure develops from the field of small density fluctuations as in standard gravitational collapse.  In order to construct a reasonable MOND cosmology which has this attribute, one must supplement the BM theory with several assumptions which may reasonably follow from a more general theory.  The first of these assumptions-- also an aspect of the earlier MOND cosmology--  is that the MOND acceleration parameter, $a_o$, which is comparable to the acceleration in the outer regions of galaxies ($\\approx 10^{-8}$ cm/s$^{-2}$), does not vary with cosmic time.  Numerically, $a_o \\approx cH_o/6$ which suggests that modified dynamics may reflect the influence of cosmology on local particle dynamics. If $a_o$ varies as the Hubble parameter, then the argumentation presented here would be incorrect.  However, it is also possible that $a_o$ is related to the cosmological constant (Milgrom 1999) and is independent of cosmic time.  If this is true then MOND plays no role in the evolution of the early radiation-dominated Universe since cosmic deceleration greatly exceeds $a_o$ on relevant scales (e.g. the Jeans length).  In the later matter-dominated, pressureless evolution, the cosmic deceleration on co-moving scales corresponding to galaxies or clusters falls below $a_o$ and one might expect modified dynamics to affect the formation of such structure.  The second assumption directly concerns the problem of the zeroth order Hubble flow; we wish to construct a theory in which MOND plays no role in the absence of fluctuations, and the background cosmology is essentially unaltered.  In other words, MOND should apply only to peculiar accelerations-- the accelerations developing from inhomogeneities-- and not to the overall Hubble flow; i.e., no MOND in a homogeneous Universe.  This assumption can find some justification in the context of a stratified scalar-tensor theory in which MOND phenomenology results from a scalar force that becomes dominant in the limit of low scalar field gradients (Sanders 1998).  The third assumption concerns the influence of the Hubble flow on the internal dynamics of an otherwise isolated spherical region. In modified dynamics, and any covariant extention thereof, it must be the case that the internal dynamics of a sub-system is influenced by the presence of an external field-- the ``external field effect\" (Milgrom 1983).  This is essentially an observational requirement on MOND imposed by the absence of discrepancies in Galactic star clusters. In other words, the underlying theory should not respect the equivalence principle in its strong version.  With respect to cosmology, it is not clear how the external field effect would come into play, but I assume here that, for an over- or under-dense spherical region, the deceleration or acceleration of the Hubble flow is the one and only external field which influences the development of the inhomogeneity.  Because the de-acceleration of the Hubble flow increases linearly with scale, fluctuations on small comoving scales are affected by MOND earlier than those on larger scale. One might expect this to lead a hierarchical scheme of structure formation, with smaller objects forming first. ",
        "conclusions": "Here I show that MOND provides the possibility of overcoming the problem of structure growth in a low-density baryonic Universe. In the context of the simple two-field non-relativistic theory of modified dynamics presented here, we see that when the background deceleration of the Hubble flow over a given scale falls below the critical MOND acceleration, $a_o$, then the growth of structure on that scale is greatly enhanced relative to the Newtonian expectation.  The growth of over-densities on smaller scale is even more enhanced due to the fact that smaller regions enter the MOND regime earlier;  the early growth of small-scale fluctuations can compensate for the effect of Silk damping on these scales.  Thus the resulting power spectrum, apart from the oscillations, closely resembles that of the favored $\\Lambda$CDM cosmology.  On the scale of galaxies (1.5 Mpc), even though the typical initial over-density is on the order of $2\\times 10^{-10}$, the fluctuation grows to the non-linear regime by a redshift of 2.5.   Thus MOND would appear not only to explain the observed large scale structure, but also provide a mechanism for early galaxy formation.  This is all achieved with the value of $a_o$ determined from galaxy rotation curves. The minimalist MOND theory has not been fine-tuned in any sense to match the observed power spectrum; the single adjustable parameter $\\beta$ lies within the range which is consistent with the observed form of galaxy rotation curves.  Of course, these conclusions depend upon the approximate validity of the assumed theory described in section 2. This theory, and assumptions embodied therein, guarantee that the early, radiation-dominated evolution of the universe is identical to that of the standard model, that the basic Hubble flow is unaffected by MOND, and that fluctuations on the scale of galaxies to super-clusters enter the MOND regime, determined by the background Hubble flow, sufficiently early (but not too early) to assure growth to the present amplitude.  Because of the necessity of such an {\\it ad hoc} theory in the absence of a more fundamental covariant theory, it is perhaps premature to compare in detail the predicted power spectrum or peculiar velocities with observations. In particular, the oscillations (Fig.\\ 2) may not actually be evident in the evolved power spectrum due to non-linear aspects of the theory which are ignored here-- specifically, not only individual Fourier components but also over-dense spherical regions cannot be considered in isolation (larger scale peculiar accelerations contribute to the background field).  The detailed results shown in Figs.\\ 1 and 2 should be taken as a demonstration that MOND, in a low-density baryonic universe, can provide a vigorous growth of fluctuations-- growth which is sufficiently rapid to lead to the large scale structure observed at the current epoch.  Finally, I re-emphasize that the presence of a dynamically significant cosmological constant plays a necessary role in the rapid growth of structure with this version of modified dynamics.  The MOND growth of inhomogeneities accelerates at the epoch when $g_b\\approx 0$ due to the dominance of the non-linear term in eq.\\ 28 (only possible with a cosmological constant comparable to ${H_o}^2$). This adds a new aspect to an anthropic argument originally given by Milgrom (1989):  we are observing the Universe at an epoch when $\\Omega_\\Lambda$ has only recently emerged as the dominant term in the Friedmann equation because it is only then that structure formation proceeds rapidly.  If the evidence in support of a baryonic-$\\Lambda$ universe continues with further observations of the CMB angular power spectrum, then some unconventional mechanism for the formation of structure must be invoked.  Here it is evident that modified dynamics, with a well-documented success in explaining the kinematic observations of galaxies and clusters without dark matter, may also successfully address the problem of structure formation in a low-density baryonic universe.  I am grateful to Art Wolfe, Eric Gawiser and Kim Griest for teaching me all about anisotropies in the CMB and the standard treatment of the evolution of density fluctuations. I thank Stacy McGaugh for stimulating this work, Arthur Kosowsky and Jacob Bekenstein for helpful criticisms, and the referee, Jim Peebles, whose numerous critical remarks led to a considerable improvement in the content and presentation of this paper.  Finally I am most grateful to Moti Milgrom for his typically penetrating comments on possible MOND cosmologies.  \\newpage"
    },
    "0011/astro-ph0011462_arXiv.txt": {
        "abstract": "In this paper, the cosmic microwave background (CMB) anisotropy in a multiply-connected compact flat 3-torus model with the cosmological constant is investigated. Using the COBE-DMR 4-year data, a full Bayesian analysis revealed that the constraint on the topology of the flat 3-torus model with low-matter-density is less stringent. As in compact hyperbolic models, the large-angle temperature fluctuations can be produced as the gravitational potential decays at the $\\Lambda$-dominant epoch well after the last scattering. The maximum allowed number $N$ of images of the cell (fundamental domain) within the observable region at present is approximately 49 for $\\Omega_m=0.1$ and $\\Omega_\\Lambda=0.9$ whereas $N\\sim8$ for $\\Omega_m=1.0$ and $\\Omega_\\Lambda=0$. ",
        "introduction": "For a long time, cosmologists have assumed the simply connectivity of the spatial hypersurface of the universe. If it is the case, the topology of closed 3-spaces is limited to that of a 3-sphere if Poincar\\'{e}'s conjecture is correct. However, if we assume that the spatial hypersurface is multiply connected, then the geometry of spatially finite models can be flat or hyperbolic as well. The metric describes only the local geometry. We should be able to observe the imprint of the ``finiteness'' of the spatial geometry if it is multiply connected on scales of the order of the particle horizon or less, in other words, if we live in a ``small universe''. \\\\ \\indent For flat multiply connected models without the cosmological constant, various constraints using the COBE DMR data have been obtained \\cite{Sokolov,Starobinsky,Stevens,Oliveira1,Oliveira2,Levin1,Roukema}. Assuming that the initial power spectrum is scale-invariant ($n\\!=\\!1$) the suppression of the temperature fluctuations on scales beyond the typical size of the cell $L$\\footnote{$L$ can be defined as twice the diameter which is defined as the maximum of the minimum geodesic distance between two points over the space.} leads to a decrease in the large-angle power. For 3-torus models without the cosmological constant in which the cell (fundamental domain) is a cube, the constraint is $L>0.8 R_\\ast$ where $R_\\ast$ is the comoving radius of the last scattering surface(horizon radius) \\cite{Sokolov,Starobinsky,Stevens,Oliveira1}. It should be emphasised that the constraint itself \\ti{does not imply that the ``small universe'' is ruled out} since the possible maximum expected number of the copies of the cell within the observable region at present is approximately 8. For models in which one side of the cell is longer than the others, the constraint for the smallest topological identification scale(=the diameter of the smallest ball which can wrap around the space) can be less stringent\\cite{Roukema}. \\\\ \\indent On the other hand, recent observations of distant supernova Ia \\cite{Perlmutter,Riess} imply the existence of ``missing energy `` which possesses a negative pressure and equation-of-state ($w\\equiv p/\\rho$) in the form of either a cosmological constant, vacuum energy density or a slowly-varying spatially inhomogeneous component ``quintessence''. Together with the first observation of the height and the position of the first Doppler peak by {{\\sc Boomerang}}\\cite{Boomerang} and {{\\sc MAXIMA}} \\cite{MAXIMA} and the COBE-DMR data, the constraint for the cosmological constant is $0.69<\\Omega_\\Lambda<0.82$ and for the matter(cold dark matter, baryon plus radiation) $0.28<\\Omega_m<0.42$ assuming adiabatic initial perturbations\\cite{Balbi}. \\\\ \\indent For low-matter-density models with flat geometry a bulk of large-angle cosmic microwave background(CMB) fluctuations can be produced as the gravitational potential decays at the $\\Lambda$-dominant epoch $1+z\\sim (\\Omega_\\Lambda/\\Omega_0)^{1/3}$. Recent works \\cite{Aurich,Inoue2,CS,AS} have shown that the angular power spectrum $C_{l}$ is completely consistent with the COBE-DMR data for some compact hyperbolic models which are incompatible with the previous analyses \\cite{Bond1}. Because the angular sizes of fluctuations produced at the late epoch are large compared to those on the last scattering for flat or hyperbolic geometry, we expect that the constraints for compact flat models with low-matter-density can be also significantly loosened. \\\\ \\indent In this paper, the CMB anisotropy in a multiply-connected compact flat 3-torus model with or without the cosmological constant is investigated. In sec II we briefly describe the time evolution of the scalar perturbation in the locally flat Friedmann-Robertson-Walker models which will be used for computing the CMB anisotropy. In sec III we compare the angular power spectrum in a flat 3-torus model with the cosmological constant to the one in the ``standard'' 3-torus model with $\\Omega_{tot}\\!=\\!1$. In sec IV a full Bayesian analysis using the COBE-DMR data has been carried out for giving the constraint on the minimum size of the cell. \\\\ \\indent ",
        "conclusions": "In this paper the CMB anisotropy in a flat 3-torus model with or without the cosmological constant has been investigated. Using the COBE-DMR data, we have done a full Bayesian analysis incorporating the effect of anisotropic correlation. It has turned out that the constraint on the low-matter-density model is less stringent compared to the ``standard'' model with $\\Omega_m\\!=\\!1.0$. The reason is that the large-angle fluctuations (on COBE scales) are produced at late time well after the last scattering. The physical size of these fluctuations is of the order of the size of the cell or less. Hence the effect of the non-trivial topology becomes insignificant. We expect that the result does not significantly change for other compact flat models with different topology\\cite{Levin1} or with different shape of the cell since the physical effect does not change. \\\\ \\indent We have seen that the analysis using the angular power spectrum is not enough since the background geometry is globally anisotropic. Even if the power spectrum well agrees with the data, it does not guarantee the validity of the model since the power spectrum has information of only isotropic components in the 2-point correlations. If the background geometry is globally anisotropic then the fluctuations form an anisotropic Gaussian field for a given orientation. On the other hand, the fluctuations become non-Gaussian if one marginalises the distribution function over the orientation\\cite{Magueijo}. For locally Friedmann-Robertson-Walker models which are spatially compact, we expect to observe zero skewness but non-zero kurtosis provided that the initial fluctuation is Gaussian. \\\\ \\indent However, for low-matter density models, anisotropy in the statistically averaged correlation on large angular scale cannot be so large since the physical size of observed temperature fluctuations that has been prodeced well after the last scatteing is much smaller than the actual size of the space(i.e. topological identification scale). \\\\ \\indent In the case of flat topology, the discrete eigenmodes have ``regular'' features which lead to significant correlations in $a_{lm}$'s. Even if marginalised over the orientation, the correlations do not completely disappear. This contrasts with the case of hyperbolic topology in which the discrete eigenmodes are ``chaotic'' and the correlations in  $a_{lm}$'s disappear if one takes an average over the position of the observer. \\\\ \\indent In order that we would be able to observe the periodic structure in the fluctuations for the 3-torus models, we need to have  $L/(2 R_\\ast)<1$ where $R_\\ast$ is the comoving radius of the last scattering surface. Because the obtained constraints give $L/(2 R_\\ast)>0.40,0.22$ for $(\\Omega_m,\\Omega_\\Lambda)=(1.0,0) ,(0.1,0.9)$, respectively, we still have a great chance of the first detection of the non-trivial topology of the universe by the future satellite missions such as MAP and \\ti{Planck} which can survey the CMB in the full sky with high resolution by measuring some specific signatures (e.g. the circles test \\cite{Cornish2} or the non-Gaussianity \\cite{Inoue3})."
    },
    "0011/astro-ph0011148_arXiv.txt": {
        "abstract": "We study static neutron stars with poloidal magnetic fields and a simple class of electric current distributions consistent with the requirement of stationarity. For this class of electric current distributions, we find that magnetic fields are too large for static configurations to exist when the magnetic force pushes a sufficient amount of mass off-center that the gravitational force points outward near the origin in the equatorial plane.  (In our coordinates an outward gravitational force corresponds to $\\partial\\ln g_{tt}/\\partial r>0$, where $t$ and $r$ are respectively time and radial coordinates and $g_{tt}$ is coefficient of $dt^2$ in the line element.)  For the equations of state (EOSs) employed in previous work, we obtain configurations of higher mass than had been reported; we also present results with more recent EOSs.  For all EOSs studied, we find that the maximum mass among these static configurations with magnetic fields is noticeably larger than the maximum mass attainable by uniform rotation, and that for fixed values of baryon number the maximum mass configurations are all characterized by an off-center density maximum. ",
        "introduction": "\\label{sec:intro}  Over the years, the typical magnitudes of the surface magnetic fields of pulsars---as inferred from measured spindown rates and simple magnetic dipole models---have been $\\sim 10^{12}-10^{13}$ G \\citep{tayl93}. Assuming magnetic flux conservation, fields of $\\sim 10^{12}$ G would naturally arise from typical main sequence star surface magnetic fields of $\\sim 100$ G during a decrease in radius by a factor of $\\sim 10^5$ \\citep{shap83}. At the extreme end of fields attainable by flux conservation, the largerst observed white dwarf magnetic field of $5\\times 10^8$ G leads to a neutron star field of $\\sim 1\\times 10^{14}$ G \\citep{carr96}, while the largest observed main sequence stellar magnetic field of $3.4 \\times 10^4$~G \\citep{bohm89} also suggests a possible field of a few times $10^{14}$~G.  Several independent circumstantial arguments link the class of objects known as ``soft $\\gamma$-ray repeaters'' (SGRs), and perhaps the so-called ``anomalous X-ray pulsars'' (AXPs), with neutron stars having magnetic fields $\\gtrsim 10^{14}$ G---the so-called ``magnetars'' \\citep{dunc92,usov92,pacz92,thom95,thom96,vasi97a}. (Table \\ref{tbl-0} displays some observed properties of these objects.) In addition to the circumstantial arguments, more direct evidence for magnetic fields of $2-8\\times 10^{14}$ G is available for two of the five known SGRs, in the form of measured periods and spindown rates of associated X-ray pulsars \\citep{kouv98,kouv99}.\\footnote{For ongoing discussion of this interpretation of X-ray timing data see e.g. \\cite{mars99a,wood99a,hard99,mars99b,chat00}.} Furthermore, the observed X-ray luminosities of the AXPs may require a field strength $B\\gtrsim 10^{16}$ G \\citep{chat00,heyl98b}. The population statistics of SGRs suggest that magnetars may constitute a significant fraction ($\\gtrsim 10\\%$) of the neutron star population \\citep{kouv94,kouv98}. As mentioned above, there are isolated examples of progenitor stars which could yield fields of $\\sim 10^{14}$ G by flux conservation, but these isolated examples do not seem sufficient to account for a significant fraction of the neutron star population. Thus, it seems likely that some mechanism {\\em generates} magnetic fields in nascent neutron stars. For example, \\cite{dunc92} suggest that the smoothing out of differential rotation and convection could generate fields as large as $3\\times 10^{17} (P_i/1\\ \\rm{ms})^{-1}$ G, where $P_i$ is the initial rotation period of the neutron star.  These considerations motivate study of the effects of ultra-strong magnetic fields on neutron star properties. In this we have been inspired by the pioneering work of Bocquet, Bonazzola, Gourgoulhon \\& Novak (1995), who performed relativistic calculations of axisymmetric neutron star structure in which the standard stress-energy tensors of a perfect fluid and the electromagnetic field were employed, and were comparable in magnitude. The maximum fields they found were of order $10^{18}$ G, with increases of 13 to 29\\% in the maximum mass of nonrotating stars for various equations of state.  An additional motivation is provided by the recent findings that magnetic fields of strengths larger than $10^{16}$ G affect the EOS of dense matter directly through drastic changes in the composition of matter \\citep{cha96,cha97,yz99,bpl00}.  The EOS is altered by both the Landau quantization of the charged particles (such as protons, electrons, etc.) and the interactions of the magnetic moments, including the anomalous magentic moments of the neutral particles (such as the neutron, strangeness-bearing $\\Lambda$-hyperon etc.) with the magnetic field.  In this work, we consider only the effects of the magnetic field on the structure, through its influence on the metric, in order to facilitate a comparison with the earlier work of \\cite{bocq95}.  The additional changes caused by the direct effects of the magnetic field on the EOS will be reported in a future work \\citep{card00}.  While \\cite{bocq95} also presented solutions for rotating neutron stars endowed with large magnetic fields, in this work we present only static solutions. In terms of the potential observability of the effects of large magnetic fields, the most relevant situation appears to be for non-rotating stars.  Neutron stars with the highest inferred magnetic fields---the so-called ``magnetars''---are all observed to be rotating very slowly, with periods on the order of seconds.  The effects of such slow rotation should have a negligible impact on the neutron star structure.  In this paper we extend the work of \\cite{bocq95}.  The theoretical formalism is outlined in \\S\\ref{sec:form}, which serves to put the problem in context.  In \\S\\ref{sec:himass} we shed light on an issue left somewhat unclear by \\cite{bocq95}: What physically determines the maximum mass and magnetic field for a given maximum density or given baryon mass?  In order to explore these questions we have chosen to solve the structural equations using a Green's function technique rather than the spectral technique employed by \\cite{bocq95}, and we also searched for the maximum mass, for a given magnetic field distribution, in a different way.  An appendix describes our numerical methods and tests of our code.  In \\S\\ref{sec:constmb} we present an illuminating view of constant baryon mass and constant magnetic moment sequences, and present higher mass configurations than those found by \\cite{bocq95} for the equations of state (EOSs) they employed.  In addition, we present the results of analogous calculations using three more recent EOSs.  Summary and outlook are contained in \\S\\ref{sec:sumout}. ",
        "conclusions": "\\label{sec:sumout}  In summary, we present a method of computing the structure of axisymmteric relativistic stars that combines elements of previous approaches, and report tests of our code. A quantitative method of determining the outer envelope (in the mass vs. radius plane) of configurations attainable with poloidal magnetic fields governed by a constant ``current function'' [see equation (\\ref{current})] has been found: magnetic fields are too large for static configurations to exist when the magnetic force pushes a sufficient amount of mass off-center that the gravitational force points outward near the origin in the equatorial plane. (In our coordinates an outward gravitational force corresponds to $-\\partial \\nu/\\partial r>0$.) We obtain larger masses of neutron stars in ultra-strong magnetic fields than have been reported previously for various equations of state, and performed computations with three representative modern EOSs. Sequences of constant baryon mass and constant magnetic moment are displayed. For all EOSs studied, the maximum attainable mass of static stars with a magnetic field determined by a constant current function is noticeably larger than that attainable with uniform rotation and no magnetic field.  The results presented here are only an initial step in exploring possible configurations of neutron stars with strong magnetic fields. As we mention below, configurations with azimuthal field components will be of physical interest, which implies that three-dimensional geometries should also be considered. But even with attention restricted to poloidal fields, we have only scratched the surface of possible configurations, as we have only considered a single current function. \\cite{bocq95} make brief mention of computations with a few other current functions.  We have performed a handful of exploratory computations using a polytropic EOS and other current functions and have found some toroidal solutions.  These toroidal solutions were not attainable with the computational approach of \\cite{bocq95}, since their method involved the specification of a finite density at the center of the star. In the case of toroidal configurations, the simple condition determining the boundary of existence of static configurations will have to be generalized, since there is no matter at the center.  These explorations will be reported in detail elsewhere.  Our work here has focused on the effects magnetic fields have on general relativistic structure, ignoring the effects of intense magnetic fields on the EOSs.  Recently, the direct effects of magnetic fields on the EOS have also been investigated \\citep{cha96,cha97,yz99,bpl00}.  Substantial effects on the EOS above nuclear saturation densities are generally produced by fields in excess of $10^{18}$ G, which is of the order of the maximum central field strengths found in this paper.  The generic effects on the EOS include softening due to Landau quantization, which is, however, overwhelmed by stiffening due to the incorporation of the magnetic moments of the various particles in neutron star matter \\citep{bpl00}. (Note that the important $B^2/8\\pi$ term is already included in our study.) Work is in progress (Cardall et al. 2000, in preparation) to provide fully self-consistent calculations of neutron star structure, in which the direct effects of magnetic fields on the EOS will be included in addition to the structural effects considered in this work.  An issue which we defer to future work is the question of stability. For systems governed by a finite number of parameters, a generalization \\citep{sork82} of the familiar one-dimensional turning point method can be employed; see \\cite{frie88} for an application to uniformly rotating relativistic stars. However, as with differentially rotating stars, this generalized turning point method is not really applicable in the present case. This is because the need to specify a current function (or rotation law in the case of differential rotation) means that defining a particular configuration requires the specification of an {\\em infinite} number of parameters.   Another issue that needs further explication before the physical relevance of the results presented here can be fully assessed is the generation of magnetic fields. We have mentioned the mechanism of \\cite{dunc92}, the generation of magnetic fields during the smoothing of differential rotation. However, the azimuthal dragging of field lines by differential rotation leads to nonvanishing azimuthal field components, in constrast to the poloidal fields studied here. It is not clear whether fields with azimuthal components would evolve into poloidal configurations, or whether there are mechanisms to directly generate strong poloidal fields. It would be of interest to explore the possibility of finding stationary solutions with toroidal magnetic fields, and in three dimensions. While this would involve more nonzero metric components, perhaps methods similar to those employed in this paper could be employed; see \\cite{bona98}.   We wish to thank E. Gourgoulhon for helpful communications concerning the calculations of \\cite{bocq95}. We are grateful to Dany Page and Ralph Wijers for their help in the preparation of Table 1. Research support from DOE grants FG02-87ER40317 (for CYC and JML) and FG02-88ER-40388 (for MP) are gratefully acknowledged.   \\appendix"
    },
    "0011/gr-qc0011049_arXiv.txt": {
        "abstract": "{ We discuss  scenarios in which  the galactic dark matter in spiral galaxies is described by a long range coherent field which settles in a stationary configuration that might account for the features of the galactic rotation curves. The simplest possibility is to consider  scalar fields, so we discuss in particular, two mechanisms that would account for the settlement of the scalar field in a non-trivial configuration in the absence of a direct coupling of the field with ordinary matter: topological defects, and spontaneous scalarization.}  \\vskip 1.5cm \\noindent PACS number(s): 11.27.+d, 04.40.-b, 98.62.Gq  \\newpage \\newpage ",
        "introduction": "It has been known for a long time that the motion of the stars and gases around the center of most galaxies can not be explained in terms of the luminous matter content of the galaxies, at least not within the context of Newtonian gravity (see \\cite{rubin} for a review). The standard view is that there is in almost every galaxy a large component of non-luminous matter (the gravitational dark matter) that forms an halo around the galaxy and that provides the additional gravitational attraction needed to explain the ``rotation curves'' in terms of standard gravitational theory. There are several proposals for this  dark component, ranging from new exotic particles such as those predicted by supersymmetry \\cite{turner,primack,susy}, to other less exotic candidates such as massive neutrinos, all collectively known as  WIMPs (weakly interacting massive particles) (see \\cite{turner,kolb} for a review), to the relative mundane idea of dark but ordinary bodies such as Jupiter-like objects collectively known as MACHOs (Massive Compact Halo Objects)\\cite{carr}. Searches for these types of objects have been made \\cite{alcock}, and although they report some findings, there doesn't seem to be enough of these objects to account for the galactic dynamics. Moreover there are severe bounds on the amount of  baryonic matter in the universe arising from big bang nucleosynthesis and for some values of the Hubble constant those bounds also imply that some of the galactic dark matter ought to be exotic\\cite{peebles1,weinberg,copi}. Independently of this and despite their popularity, these type of models suffer from various problems and require surprising coincidences (see for example \\cite{sellwood}).\\\\ \\\\ Another type of proposal, which is in some sense more radical, is based on the idea that the gravitational theory would have to be modified when dealing with the scales associated with the motion of stars in galaxies \\cite{finzi,sanders}, in particular the idea is embodied by the proposal of Milgrom \\cite{milgrom}, that the laws of motion are modified  when the accelerations involved are extremely small. Unfortunately this scenario has not, as yet, been converted into a fully relativistically invariant theory. Another type of model that has been exposed is to replace general relativity by a  higher order in curvature theory, which in  some particular cases appears to be obtaining encouraging results\\cite{Manheim}. The problem with this approach is that these types of theories have in general problems of principle like for example the lack of a well posed initial value formulation. Nevertheless such relatively radical proposals are still attractive, due in part to an intrinsic problems of the more conservative approaches in explaining the generality and universality of the phenomena, namely the   fact that the amount of luminous matter seems to be such a good indicator of the amount of the dark matter component \\cite{persic} and the fact that the dark component happens always to distribute itself in such a way that the resulting rotation curves (hereafter referred to as RC) are almost flat away from the  galactic centers \\cite{rubin1}.  \\\\ \\\\ Thus, in contrast with the  former scenarios which would need to assume not only the existence of the dark matter but also give some evolutionary scenarios that result in the aforementioned universality in its distribution, the modified gravity scenario would naturally account for such correlations without the need for additional assumptions. On the other hand the former scenarios do not present any problem in lending themselves to an acceptable theoretical formulation, compatible with present theories of particle physics and general relativity.  \\\\ \\\\ The object of this article is to discuss a third type of scenario which has some of the advantages of each scheme. The idea is to take dark matter to be described not by a bunch of particles whose distribution needs to be explained but by a coherent field which would settle in a universal stationary configuration that would account for the generic features of the RC. The simplest possibility is provided by scalar fields, which would of course have to be very long ranged (i.e. masses smaller than ${1 \\over R_G}$ where $R_G$ is the radius of the largest galaxy with flat RC). The basic problem is that there are very severe experimental bounds for the direct coupling of such a field with ordinary matter \\cite{panov}, and in the absence of such coupling the field will in general settle globally in the minimum of the potential leading to a homogeneous configuration that will not produce the desired effects. On the other hand, one could hope that, given the likelihood of existence of large black holes at the center of most galaxies, they would account for the nontrivial configuration of the scalar fields. Unfortunately these kind of situation is largely forbidden by the ``Black Hole No hair theorems'' for scalar fields \\cite{MH,DS,JDB}. These limitations severely reduce the types of models one can consider, in particular, there are, known to these authors, only three mechanisms that would account for the settlement of the scalar field in a non-trivial configuration in the absence of a direct coupling of the field with ordinary matter (or some other exotic matter which we will not consider because of the incremental number of hypothesis it involves): a)  boson-star like clumps, b) spontaneous scalarization, and  c) topological defects. Other models that lack these features have been considered, for example in \\cite{tonatiuh}. However, such models face two problems: first, they give rise to spherical configurations where the scalar field in consideration is singular ``at the center'', and second, the resulting scalar field potential needed to account for the flat RC depends explicitly on the value of the ``tangential velocity'' of stars at the flat region. That is to say, such a potential have to be adjusted differently for different galaxies. Needless to say that both problems clearly make those schemes unsuitable as models for the problem at hand. \\bigskip  Concerning the case ``a)'' mentioned above, it has been analyzed in \\cite{peebles}. Their analysis focuses on cosmological and evolutionary considerations as well as the issues related to the conditions under which the assumption of long range coherency of the scalar field is justified, rather than the universal features of the galactic rotation curves. We will deal here with the other two cases ``b)'' and ``c)''.  \\\\ \\\\ The scenario ``b)'', namely the spontaneous scalarization (see Sec.V)\\cite{Damour}, is in some sense simpler because it involves a single scalar field in contrast of the various fields needed in the simplest versions of topological defects (e.g., global monopoles). Here the mechanism that allows for the nontrivial stationary configuration of the scalar field is connected  to a non-minimal coupling of the scalar field to the curvature. This results in the effective gravitational coupling becoming dependent on the scalar field. The point is that such a coupling allows for the reduction of the total energy of the configuration (in comparison with the corresponding configuration with the same baryon number and no scalar field) for which the scalar field deviates from the trivial configuration by taking values that reduce the gravitational coupling in the regions of high matter density \\cite{SSN}. Thus the model must incorporate from the onset the non-minimal coupling that seems to be needed to account at the same time for the correlations in the dark-luminous matter components (see \\cite{NSS} and the discussion of the third scenario below). The disadvantage of this model, which is in fact shared by the first model (i.e., boson stars) \\cite{Igor}, is precisely the lack of resilience against black holes whose existence in most galaxies, if confirmed, would seem to preclude, through the no-hair theorems \\cite{MH,DS,JDB}, the models based on this mechanism.  \\\\ \\\\ The scenario ``c)'' is exemplified by the model of global  monopoles \\cite{vilenkin1} which have the notable feature of naturally leading to a $1/r^2$ energy density behavior which would naively account for the flat rotation curves and which upon taking the symmetry braking scale to be the GUT scale would result in the correct order of magnitude for the galactic dark matter. Unfortunately, upon further examination of the simplest model severe problems arise, in particular the monopole configuration turns out to be repulsive \\cite{harari}, and moreover the configuration would be too universal in the sense that it would be independent of the size of the galaxy thus defeating the hope for the correlation  of dark to luminous matter over a range of galactic sizes. There is nevertheless hope to overcome these problems by the consideration of slightly more complicated models \\cite{NSS}. In that work  the simple monopole model was supplemented by the introduction of a non-minimal coupling between the scalar fields and curvature (see \\cite{NSS} and Sec.VI). This resulted in the restoration of gravitational attraction leading to regions of relatively flat rotation curves and to the possibilities of the dark-luminous matter correlations arising from the fact that in these models the scalar potential $V(\\Phi^a \\Phi_a)$ (where $\\Phi^a$ stands for a triplet of scalar fields that characterize the global monopole) is replaced with the effective scalar potential $V(\\Phi^a\\Phi_a)+ F(\\Phi^a\\Phi_a,R)$ (here $R$ stands for the scalar curvature of the space-time metric) whose minima depend on the amount of matter present trough  the effect of the latter on $R$.  The global monopole model has the additional advantage of resilience against the formation of black holes in the galactic centers, since their topological charge makes them immune to the devastating limitations imposed by the no hair theorems. \\\\ \\\\ Despite the promising features of the model ``c)'', our intention in the present work is to take a step backwards and look at the problem from a more general point of view before embarking in the methodical study of a particular type of model. \\\\ \\\\ The article is organized as follows: in Section II, we analyze the generic form of the rotation curves of galaxies in a general relativistic context. In section III, we comment on the  Newtonian approximation and on the embedding of the galaxy in the large scale space-time. In Section IV, we discuss the additional information that can be obtained about the metric from other considerations, specifically the deflection of light by the galaxy. Section V reviews the spontaneous scalarization scenarios. In Section VI, we review the non-minimally coupled global monopole model and discuss its shortcomings. Finally, in Section VII we offer a discussion and analyze the directions for further developments. \\bigskip ",
        "conclusions": "The galactic rotation curves continue to pose a challenge to present day physics as one would want to understand not only the nature of the dark matter that is associated with them but also the reason behind their universality (i.e., why is it distributed within a galaxy in a way that leads to almost flat rotation curves ?, and why is the amount of dark matter present in a galaxy so well correlated with the luminous matter ? \\cite{persic,tully}).  Models based in ordinary physical objects could already be facing problems (depending on the exact value of the Hubble constant \\cite{copi}) in view of the bounds that big bang nucleosynthesis impose on the baryon content of the universe.  Models based on particle physics are the most commonly considered (usually within a Newtonian scheme) but they need to address the nature and the distribution problems separately, leading to a larger number of hypothesis and surprinsing coincidences \\cite{sellwood}.  In view of the recent cosmological measurements and the theories that have been put forward to explain them \\cite{5essence}, one is naturally lead to consider alternative models based in the introduction of long range coherent fields \\cite{peebles}. In this work, we have given a review of various types of approaches to these questions indicating in each case the problems and advantages.  We have argued that so far the most promising and simple approach would involve global monopoles with some sort of nonminimal coupling to gravity. This remains for the future to establish how far can this sort of ideas be pushed towards the goal of making a realistic and compelling model for the dynamics and evolution of galaxies. In particular any such model must also be studied in the context of cosmological perturbations, large scale structure and the CMB. In this regard we should point out that the simplest models of topological defects as seeds for structure formation seem to be incompatible with the acoustic peak in the CMB anisotropies detected by Boomerang and Maxima \\cite{boommax}. However, all these studies have considered the simplest minimally coupled models and it is unclear how would the models of the type being analyzed here behave in this respect. Finally, we should mention that the currently favored cosmological scenarios require at least two hypothetical components: the Cold Dark Matter (usually in the form of WIMPS) necessary for the structure growth and the dark matter in galaxies and clusters, and the cosmological constant $\\Lambda$  which provides the closure density (as required by inflation) as well as the repulsive component that seems to be required in order to account for the observations of the luminosity-distance of high red shift (type Ia) supernovae \\cite{snI}. The fact that the non-minimally coupled monopoles exhibit both an attractive regime at short distances and a repulsive regime at large distances leads us to speculate whether these type of models can be used to explain the two aspects of the unobserved energy content of the universe in terms of a single hypothetical component. Needless is to say that all these aspects will require intense further exploration, which we hope to undertake in the near future.   \\vskip 1cm {\\bf Acknowledgments}  \\bigskip U.N. is supported by a CONACyT postdoctoral fellowship grant 990490; M.S. and D.S. acknowledge partial support from DGAPA-UNAM Project No. IN121298 and from CONACyT Projects 32551-E and 32272-E. Authors thank the supercomputing department of DGSCA-UNAM.  \\newpage"
    },
    "0011/astro-ph0011524_arXiv.txt": {
        "abstract": "I give a brief overview of cosmic ray physics, highlighting some key questions and how they will be addressed by new experiments. ",
        "introduction": "There are several new experiments in cosmic-ray physics and related fields running now or planned for the near future.  These include measurements of antiprotons with balloons and from space, new gamma-ray detectors to cover the full energy range from $<100$~MeV to $>10$~TeV, studies of the knee region with a variety of experiments, giant air shower detectors for the highest energy particles, and huge, deep neutrino detectors to find astrophysical sources of high energy neutrinos.  In this talk, I will review the key questions driving these experiments and, where possible, put them in a larger context and relate them to each other.  After some brief remarks about the sun as a cosmic-ray source, I review cosmic rays of galactic origin and then extragalactic cosmic rays.  Energy content of the various components of the cosmic radiation and the corresponding power to maintain the observed fluxes are used as a guide to possible sources. ",
        "conclusions": ""
    },
    "0011/astro-ph0011238_arXiv.txt": {
        "abstract": "We present a technique that combines Zeeman Doppler imaging (ZDI) principles with a potential field mapping prescription in order to gain more information about the surface field topology of rapid rotators. This technique is an improvement on standard ZDI, which can sometimes suffer from the suppression of one vector component due to the effects of stellar inclination, poor phase coverage or lack of flux from dark areas on the surface. Defining a relationship beween the different vector components allows information from one component to compensate for reduced information in another. We present simulations demonstrating the capability of this technique and discuss its prospects. ",
        "introduction": "The solar-stellar analogy is commonly invoked to explain stellar activity phenomena in terms of solar features such as prominences and flares \\cite{radick91}. A solar-type dynamo mechanism is also thought to operate in other late-type stars.  However it is unclear if this exact same mechanism will operate in stars covering a wide range of convection zone depths, rotation rates and masses. Through the Mt Wilson H\\&K survey, which has monitored stellar rotation periods and activity cycles using chromospheric emission variations spanning a period of over 30 years, we know that many cool stars also display regular solar-type activity cycles. However, other stars in the survey show irregular variations and some show no cyclic patterns over the time span of the observations \\cite{donahue96,baliunas95}. \\scite{saar99} and \\scite{brandenburg98} find that the type of dynamo activity appears to change with rotation rate and age in a study using data from the Mt Wilson survey as well as other photometric studies of cool stars.  Doppler imaging techniques can map $T_{\\rm eff}$ flux distributions on the surfaces of rapidly rotating cool stars \\cite{cameron92doppler,piskunov90,rice89,vogt88}. Starspot maps of cool stars taken a few rotation cycles apart allow us to determine the differential rotation rate on rapid rotators accurately \\cite{donati97doppler,barnes2000pztel,petit2000HR1099,donati2000RXJ}. As \\scite{cameron00omega} report, these measurements lend support to differential rotation models in which rapid rotators are found to have a solar-type differential rotation pattern \\cite{kitchatinov99}. Doppler maps of rapidly rotating stars show flux patterns which are very different to those seen on the Sun, with polar and/or high-latitude structure often co-existing with low-latitude flux \\cite{deluca97,strassmeier96table}. This is hard to reconcile with the solar distribution where sunspots and therefore the greatest concentrations of flux tend to be confined between $\\pm 30^{\\circ}$ latitude. Does this indicate different dynamo modes are excited in these stars ? Models by \\scite{granzer00} and \\scite{schussler96buoy} explain the presence of mid-to-high latitude structure on young, rapidly rotating late-type stars in terms of the combined effects of increased Coriolis forces and deeper convection zones working to ``pull'' the flux closer to the poles. Even though \\scite{granzer00} find that both  equatorial and polar flux can exist on T Tauri stars that have very deep convection zones, in terms of flux ``slipping'' over the pole, they find the emergence of low-latitude flux on K dwarf surfaces difficult to explain.  The technique of Zeeman Doppler imaging (ZDI) allows us to map the magnetic field distributions on the surfaces of rapid rotators using high resolution circularly polarized spectra \\cite{semel89}. Magnetic field maps of the subgiant component of the RS CVn binary, HR1099 ($P_{rot}=2.8$d), and the  K0 dwarfs, AB Dor ($P_{rot}=0.5$d), LQ Hya ($P_{rot}=1.6$d) have been presented in several papers \\cite{donati97doppler,donati99doppler,donati99lq}. These maps all show patterns for which there is no solar counterpart: strong radial and azimuthal flux covers all visible latitudes. On AB Dor, strong unidirectional azimuthal field encircling the pole is consistently recovered over a period of three years. This is very different to what is seen on the Sun, where hardly any azimuthal flux is observed at the surface and mean radial fields are much smaller than those observed in ZD maps.  While ZDI is an important tool in measuring the surface flux on stars, some questions about the authenticity of the features in these Zeeman Doppler maps remain. ZDI makes no assumptions about the field at the stellar surface; the radial, azimuthal and meridional vectors are mapped completely independently \\cite{hussain99thesis,donati97recon,brown91zdi}. For the true stellar field, however, the physics of the stellar interior and surface will determine the relationship between the field components. Additionally, poor phase coverage is found to lead to an increased amount of cross-talk between radial, meridional and azimuthal field components in ZDI maps \\cite{donati97recon}. Due to the presence of starspots, these maps are also flux-censored. Indeed \\scite{donati99doppler} and \\scite{donati97doppler} find that the low surface brightness regions on AB~Dor appear to correlate with the strongest regions of radial magnetic flux. For these reasons Zeeman Doppler (ZD) maps may not present an accurate picture of the surface field on these stars.  By mapping potential fields on the surfaces of stars, we can evaluate more physically realistic models of the surface flux distribution. While we do not necessarily expect the field to be potential locally, this assumption should be adequate on a global scale \\cite{demoulin93}. \\scite{jardine99pot} evaluate the consistency of Zeeman Doppler maps for AB Dor with a potential field using maps obtained over three years. Although they find that the potential field cannot reproduce features such as the unidirectional azimuthal flux found at high latitudes, they do find a good correlation between the ZD azimuthal map and potential field configurations produced by extrapolating from the ZD radial map.  The technique presented in this paper reconstructs potential field distributions directly from observed circularly polarized profiles. This technique allows us to produce different configurations which are more physically realistic and which provide more information about the surface topology. ",
        "conclusions": "Many limitations in ZDI arise from the sensitivity of circularly polarized profiles (Stokes V parameter) to the line-of-sight components of magnetic fields. However, until spectropolarimetric instrumentation is sufficiently advanced so that all four Stokes parameters can be observed simultaneously for cool stars this problem is unlikely to be resolved. Currently, the four Stokes parameters can only be measured in bright, highly magnetic A and B stars  \\cite{wade00stokes,wade00ap}.  The lack of physical realism in ZD maps also allows for unrealistic magnetic field distributions. By limiting ZD reconstructions to potential field distributions we can overcome some of these problems. In mapping potential fields, we introduce a mutual dependence on all three surface field vectors thereby recovering more field information. The benefits of this method are threefold: firstly, in regions where one vector is suppressed due to the inclination of the star (such as regions of low-latitude meridional field) there is still a reasonably high contribution from the other two and thus field information is reconstructed more reliably; secondly, in areas of bad phase coverage, azimuthal fields can sometimes still be recovered (as their signatures are strongest about $45^{\\circ}$ away from the previous point of observation) and the mutual dependence of the field vectors can be exploited to recover some radial and meridional field information; and finally, in regions such as starspots where the field is thought to be mainly radial but where there is little flux contribution from all three vectors, the field information in the surrounding spot areas can be used to recover more of the lost polarized signal.  In addition to these improvements, potential field configurations should allow us to test if the unidirectional band of azimuthal field seen in ZD images of AB Dor is necessary to fit the data. Such a band is incompatible with a potential field and its presence challenges our current understanding of magnetic field generation.  If datasets are available that span a sufficiently long time-span it will be possible to monitor which modes are consistently active on these stars which will give us insight into the dynamo mechanisms operating in cool stars. Extrapolating these surface potential field models to the corona will enable us to model the coronal topology of the star and thus to predict where stable prominences may form in these stars.  Finally, special care must be taken when applying this technique to rotationally broadened Stokes V spectra from cool dwarfs. The weighting scheme used will have to be tested more thoroughly in order to ensure that we are not sacrificing realism for the ``simplest images'' (in terms of information content). More realistic images may be obtained by incorporating a form of entropy which pushes the field distribution to a power law form, a possibility we are investigating further."
    },
    "0011/astro-ph0011074_arXiv.txt": {
        "abstract": "The Red-Sequence Cluster Survey (RCS) is a 100 deg$^2$ optical survey for high-redshift galaxy clusters. One of the goals of the survey is a measurement of $\\Omega_m$ and $\\sigma_8$ via the evolution of the mass spectrum of galaxy clusters. Herein we briefly describe how this will initially be done, and also demonstrate the eventual power of the RCS for this type of measurement by a qualitative analysis of the first 1/10th of the survey data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011242_arXiv.txt": {
        "abstract": "{ The narrow-line Seyfert 1 galaxy \\mrk\\ was observed for 60 ks with the \\xmm\\ observatory. The source shows a complex X-ray spectrum. The 2-10 keV spectrum can be adequately represented by a power law and broad Fe K$\\alpha$ emission. Between 0.7 and 2 keV the spectrum is harder and exhibits a flux deficit with respect to the extrapolated medium energy slope. Below 0.7 keV, however, there is a strong excess of emission. The RGS spectrum shows an edge-like feature at 0.7 keV; the energy of this feature is inconsistent with that expected for an OVII edge from a warm absorber. \\mrk\\ varies by a factor of $\\sim$ 2 in overall count rate in the EPIC and RGS instruments on a timescale of a few thousand seconds, while no significant flux changes are observed in the ultraviolet with the OM. The X--ray variability is spectrally dependent with the largest amplitude variability occurring in the 0.4-2 keV band. The spectral variability can be explained by a change in flux and slope of the medium energy continuum emission, superimposed on a less variable (or constant) low energy emission component. ",
        "introduction": "\\label{sec:introduction}  X--ray emission currently provides the best tracer of conditions close to the supermassive black hole central engines of AGN. However the origin of AGN X--ray emission remains poorly understood. Between 1 and 20 keV AGN generally have a power law spectrum modified by the effects of reflection and absorption (Pounds \\etal \\cite{pounds90}); below 1 keV many AGN also show some `soft excess' emission (Walter \\& Fink \\cite{walter93}). The majority of recent models for AGN X--ray emission explain the soft excess as emission from an optically thick accretion disk, and the power law component as Compton upscattering of soft photons by a hot corona.  \\mrk\\ is a nearby (\\(z=0.0129\\)), soft X--ray bright (0.5 - 2 keV flux of $\\sim 10^{-11}$ erg s$^{-1}$ cm$^{-2}$) and variable (factor 2 in \\(\\sim\\) 1000s, Leighly \\etal \\cite{leighly96}) narrow line Seyfert 1 galaxy observed through a relatively small Galactic column of \\(\\sim 1.8 \\times 10^{20}\\) cm\\(^{-2}\\). It is therefore an extremely good candidate for studying the dynamical and radiative processes which give rise to the X--ray emission in AGN. In this letter we present the spectrum and spectral variability of \\mrk\\ as observed by the combination of instruments on board \\xmm\\ during the performance verification (PV) observation. ",
        "conclusions": "We have presented RGS, EPIC and OM spectra and lightcurves of \\mrk\\ from the 60 ks \\xmm\\ PV phase observation. In all three EPIC cameras the spectra above 2 keV are well described by a power law with a broad Fe K$\\alpha$ emission line. Relative to the power law, the EPIC spectra show a deficit of photons between $\\sim$ 0.7 and 2 keV and excess emission at lower energies.  The RGS spectra contain no strong, narrow emission lines. Warm absorber edges are not seen at the expected energies of OVII and OVIII in the RGS spectra; an edge-like feature at 0.7 keV could only be due to OVII if the absorbing material were redshifted more than 10,000 km/s relative to the system. Significant variability is detected in the X--ray lightcurves and hardness ratios, while no variability is detected in the UV with the OM. A variable warm absorber is not required to account for the spectral variability, which can be explained by a medium energy continuum component which softens as it increases in flux, with an additional contribution from a less variable emission component at low energy. Changes in magnetic fields threading the accretion disk corona, or the spectrum of particles injected into the corona, can explain the observed spectral variability."
    },
    "0011/astro-ph0011418_arXiv.txt": {
        "abstract": "Rossi X-ray Timing Explorer All Sky Monitor observations of \\src\\ show that the source experienced an outburst in January to April 2000 reaching a peak luminosity of greater than \\sqig10$^{38}$ ergs s$^{-1}$. RXTE Proportional Counter Array observations during this outburst reveal the presence of pulsations with a 2.37s period.  However, optical photometry of the optical counterpart showed the source to be still significantly fainter than it was more than half a year after the outburst in the 1970s when \\src\\ was discovered. ",
        "introduction": "The first three X-ray point sources to be discovered in the Small Magellanic Cloud (SMC) were SMC X-1, SMC X-2 and SMC X-3, all of which are thought to be X-ray binaries. While SMC X-1 is a persistent pulsing source (e.g. Wojdowski et al. 1998 and references therein), SMC X-2 and X-3 are transient sources that were discovered in 1977 with SAS-3 (Clark et al. 1978) which were both initially very bright at a level of \\sqig 1.0 $\\times$ 10$^{38}$ ergs s$^{-1}$ (2 $-$ 11 keV) for assumed distances of 71 kpc and remained in outburst for approximately one month (Clark, Li, \\& van Paradijs 1979). Both sources were also detected with HEAO-1 A-2 experiment (Marshall et al. 1979).  However, they have since exhibited little activity.  SMC X-2 was not detected in an Einstein IPC survey (Seward \\& Mitchell 1980) and was seen in only one of two observations made with the ROSAT PSPC detector (Kahabka \\& Pietsch 1996, Haberl \\& Sasaki 2000) at a luminosity of 2.7 $\\times$ 10$^{37}$ ergs s$^{-1}$ (0.15 $-$ 2.4 keV). SMC X-3 was also not detected by Seward \\& Mitchell (1980) but was seen in one ROSAT HRI observation (Haberl \\& Sasaki 2000). An optical counterpart for SMC X-2 was proposed by Crampton, Hutchings, \\& Cowley (1978) which was resolved into a pair of stars, ``A'' and ``B'', with the fainter object (``B''), a main-sequence Be star, identified as the likely counterpart by Murdin, Morton, \\& Thomas (1979).  Infra-red observations of SMC X-2 during a quiescent period are reported by Coe et al. (1997). In none of the previous X-ray observations of SMC X-2 or SMC X-3 were any pulsations detected although they are expected to be present as most Be star X-ray sources are X-ray pulsars. Another bright X-ray pulsar candidate within the SMC is the little studied source H0107-750 which also has a Be star as its proposed optical counterpart (Whitlock \\& Lochner 1994 and references therein).  Although no other SMC X-ray pulsars apart from SMC X-1 were known for many years there has recently been an explosion in the discovery of SMC X-ray pulsars. Hughes (1994) discovered a 2.7s pulsar, and three pulsars were found to be active in a single RXTE observation of the vicinity of SMC X-3 (Corbet et al. 1998). Further RXTE, ASCA and SAX observations have now lead to the identification of over 20 X-ray pulsars in the SMC (e.g. Haberl \\& Sasaki 2000, Yokogawa et al. 2000, and references therein) and their variability and several optical identifications suggests that at least the vast majority are transient Be star systems. The number of known Be star X-ray binaries in the SMC is thus now starting to rival the number known in the Galaxy (e.g. Bildsten et al. 1997).  In this paper we report the detection of activity from the vicinity of SMC X-2 in early 2000 with the Rossi X-ray Timing Explorer All-Sky Monitor (ASM). Subsequent observations with the RXTE Proportional Counter Array (PCA) then revealed the presence of 2.37s pulsations. Contemporaneous optical photometry of the optical counterpart are compared to the initial discovery photometry. Note that in this paper we use an assumed distance of 65 kpc to the SMC for consistency with several recent publications (e.g. Kahabka \\& Pietsch 1996, Haberl \\& Sasaki 2000) which is somewhat smaller than the 71 kpc adopted by Clark et al. (1978). ",
        "conclusions": "Given the large number of X-ray pulsars in the SMC can we be certain that the new pulsar that we have detected is indeed SMC X-2? The pulse period is rather short and we can make use of the result of Stella, White \\& Rosner (1986) who noted an anti-correlation between pulse period and maximum X-ray luminosity. SMC X-2 at its peak observed with SAS-3 and the RXTE ASM was rather bright at $>$ \\sqig 10$^{38}$ ergs s$^{-1}$ which, together with the pulse period that we measure, would be consistent with the relationship found by Stella et al.  The vicinity of SMC X-2 may also be less densely populated with active pulsars than the region around SMC X-3. After emission had been detected from the vicinity of SMC X-3 with RXTE subsequent analysis revealed the presence of three pulsars (Corbet et al. 1998), while ASCA observations showed that none of these were in fact SMC X-3 itself. However, in the case of SMC X-2, we only detect one pulsar and, in addition, its location is consistent at the \\sqig 3 arc minute level with SMC X-2. The brightening observed in the ASM light curve for SMC X-2 is also consistent with the source being not more than a few arcminutes from SMC X-2. Further evidence still for the identification of the 2.37s pulsar with SMC X-2 comes from the detection of this source by Torii et al. (2000) with the Gas Imaging Spectrometer experiments on board ASCA. The GIS is an imaging instrument which gives source locations with 90\\% errors of 1 to 2 arc minutes radius (Tashiro et al. 1995). In the ASCA GIS survey of the SMC (Yokogawa et al. 2000) 12 X-ray pulsars and 8 candidate X-ray pulsars were identified. We thus obtain an approximate estimate of the chance of finding any X-ray pulsar or candidate within 2 arc minutes of SMC X-2 in a snap-shot survey of a few percent.\\ Another independent estimate of the probability that the source that we have observed is SMC X-2 comes from the serendipitous RXTE slew detection (observation \\#1). The slews across the SMC that occur during routine RXTE observations provide a monitor for bright new sources within this entire region.  There are currently only five sources known in the SMC with luminosities \\sqig10$^{38}$ ergs s$^{-1}$: SMC X-1, SMC X-2, SMC X-3, H0107-750, and the 31 second pulsar XTE J0111.2-7317 (Chakrabarty et al. 1998). The probability of finding a new source with a comparable brightness at a position consistent with one of these known bright sources within the approximately 2\\degrees\\ $\\times$ 2\\degrees\\ extent of the SMC is less than 1\\%.  Together, the comparable luminosity observed with RXTE and SAS-3 and the positional coincidence at the few arc minute level, strongly suggest that the source we have observed is indeed SMC X-2. The evidence would be strengthened further still if the position of the 2.37 s pulsar could be localized more precisely or evidence found for pulsations at this period in the archival SAS-3 or HEAO-1 observations.  The outburst we have observed with RXTE could be due to either changes in the circumstellar envelope of the Be star, associated with orbital variability, or both of these effects.  We note that no other outburst of a similar size is apparent in the ASM light curve which would argue for a short lived change in the circumstellar envelope.  If \\src\\ has properties consistent with the general correlation between orbital and pulse periods exhibited by most Be/neutron star binaries (Corbet 1986) then an orbital period of order 20 days to within a factor \\sqig2 might be expected. However, no evidence for any periodic variability is found in the ASM light curve.  During the course of our observations the pulse period of \\src\\ increased and was longer still at 2.37230 $\\pm$ 0.00004 s when ASCA observed on April 25-26 (Torii et al. 2000). This change in pulse period may be due to orbital motion and would correspond to a velocity change of \\sqig 100 km s$^{-1}$. While there are currently too few pulse period measurements to place any strong constraints on the orbital parameters, if \\src\\ returns again to an active state it is possible that additional observations may yield the orbital period, either from monitoring the pulse period or from periodic outbursts that the ASM could observe.  The pulse period that we find for \\src\\ is among the shortest currently known for high mass X-ray binaries (e.g. Bildsten et al. 1997) especially for those systems which contain Be star primaries. As noted above, this short pulse period is consistent with the high source luminosity. SMC X-3, which also exhibited a high maximum luminosity, may also be expected to have a relatively short pulse period. Previously Clark et al. (1978, 1979) argued that the high luminosities of SMC X-1, X-2 and X-3 implied that the luminosity distribution of the SMC X-ray sources is shifted toward higher luminosities compared to Galactic systems. However, now that the transient X-ray pulsar population of the SMC is known to be so large this shift to higher luminosities is called into question (see e.g. Haberl \\& Sasaki 2000) as the earlier observations were only detecting the high-luminosity tip of the large SMC X-ray pulsar population. Extensive studies of these objects down to low flux levels are required to fully establish their overall properties as a class."
    },
    "0011/astro-ph0011132_arXiv.txt": {
        "abstract": "I summarize some of the consequences for the optical and very-near-infrared spectra of T dwarfs (in particular) and brown dwarfs (in general) of their possible dominance by the neutral alkali metal lines.  As a byproduct of this study, I estimate the true optical color of ``brown'' dwarfs. ",
        "introduction": "The early discovery phase for L dwarfs and T dwarfs has ended and a major focus is now on their characterization.  The atmospheres of brown dwarfs are dominated by H$_2$, H$_2$O, CH$_4$, NH$_3$, the neutral alkali metals, and grains, but how theory translates this basic knowledge into effective temperatures, gravities, and compositions has yet to be determined.  Establishing the spectral and color diagnostics that are most appropriate for L/T studies is complicated by ambiguities in the cloud/grain models and a paucity of opacity data.  In particular, though T dwarfs are being informally defined by their methane features at 1.7 \\mic and 2.2 \\mic, the methane database itself is far from complete.  The methane opacities on the red side of the $H$ band are certainly in error by a factor of 3 to 5 (witness Gliese 229B\\cite{leg99}) and the hot bands are completely missing.  The latter means that even the sign of the opacity's dependence upon temperature can be in error. Nevertheless, there has been great overall progress towards understanding what makes these objects unique and what their spectra are telling us.  In this paper, I sidestep a comprehensive study of these issues and summarize three interesting topics in brown dwarf theory that have emerged of late.  They are 1) what determines T dwarf spectra shortward of 1.0 micron, 2) what is the true color of a ``brown\" dwarf, and 3) what is the effect of heavy element depletion (``rainout'') on the abundance profiles of the neutral alkali metal atoms. A subtext of this contribution is the central importance of the alkali metals in spectrum formation. ",
        "conclusions": "L and T dwarf spectra are unique among ``stars\" and require new databases, approaches, and thinking to fully understand.  Exploring as we are new worlds, we will require new tools and instincts with which to navigate.  Along with accurate cloud models, methane, and water, the alkali metals hold the key to unraveling the mysteries of the substellar objects that we now know inhabit the solar neighborhood in abundance.  \\subsubsection{Acknowledgments:\\\\} I thank my long-time collaborators, Jonathan Lunine, Bill Hubbard, and Mark Marley, for simulating input and both Davy Kirkpatrick and Neill Reid for an advanced glimpse at their stunning 2MASSW J1507 spectrum. This work was supported in part by NASA under grants NAG5-7073 and NAG5-7499."
    },
    "0011/astro-ph0011060_arXiv.txt": {
        "abstract": "We probe the region around the protostar HH108MMS by deep mid infrared photometric and polarimetric imaging.  The protostar is detected at 14$\\mu$m in absorption against the diffuse background.  Next to HH108MMS, we find a second absorbing core, named Q1, and the young stellar object IRAS18331--0035 which is more advanced in its evolution and already seen in emission at 12$\\mu$m and 14$\\mu$m.  HH108MMS, Q1 and IRAS18331--0035 form a triplet along an extended filamentary absorption feature.  From the variation of the surface brightness across the source, we derive for HH108MMS and Q1 the optical depth and density profile.  Along the axes which are parallel to the filament, the density distributions follow a $\\rho\\propto r^{-1.8}$ power law. We estimate that the intensity of the background radiation at 14$\\mu$m is about two times stronger than the intensity of the interstellar radiation field in the solar neighborhood. The present photometric data of IRAS18331--0035 between 12$\\mu$m and 1.3mm can be explained by a central source with a luminosity of 2.5\\,\\Lsun\\ that is surrounded by a spherical cloud of 1.1\\,\\Msun\\ with a $1/r$ density distribution. As HH108MMS is also seen in the millimeter dust emission, we can derive the ratio of the dust extinction coefficients at 14$\\mu$m and 1.3mm and obtain $\\kappa_{14\\mu\\rm m} / \\kappa_{1300\\mu \\rm m}\\sim 470$.  Because models for the dust in the diffuse interstellar medium predict a ratio of around 2000, our value points to fluffy composite grains which are expected to prevail in dense and cold environments.  \\noindent First mid infrared polarisation images of pre--stellar absorbing cores are presented.  At 12$\\mu \\rm{m}$ and 14$\\mu \\rm{m}$ the polarisation of the region around HH108MMS is strong ($\\geq 15\\%$) and tightly correlated with the source triplet.  We demonstrate that the high degree of polarisation can be explained by extinction of rotationally aligned dust particles of moderate elongation. ",
        "introduction": "\\noindent Herbig Haro objects mark shock regions in the outflow from very young stars.  The 1.3mm continuum survey of Herbig Haro objects by Reipurth et al.~(1993) indicates that the young stars, which are the energy sources of the outflow, are surrounded by dusty envelopes of about 0.1 -- 3 solar masses.  The circumstellar material is so massive that the stars are probably still in the accretion phase.  In a follow--up study of the 1.3mm dust continuum, Chini et al. (1997) discovered in the Serpens star forming region which is at a distance $D= 310$\\,pc (De Lara et al. 1991) a protostellar candidate located 71$''$ (0.11pc) away from the source IRAS18331--0035.  The latter is believed to be the driving engine of the two aligned nearby bow shocks HH108 and HH109 (Reipurth \\& Eiora 1992, Ziener \\& Eisl\\\"offel 1999); they lie to the South-West of IRAS18331--0035, about 0.1pc and 0.21pc away. In this paper, we present photometric and polarimetric mid infrared images of the region. ",
        "conclusions": "\\noindent Among many other models, the self--similarity solutions of Shu et al.~(1977) for the early evolution of protostars predict in the absence of rotation and magnetic fields a $\\rho \\propto r^{-2}$ profile in the outer envelope and a less steep distribution $\\rho \\propto r^{-3/2}$ in the inner part.  Basu \\& Mouschovias (1995), who include magnetic fields and rotation, also find a power law density distribution in the envelope with an exponent between $-1.5$ and $-1.8$ and a constant density central region of $\\sim$50\\,AU size.  Masunaga \\& Inutsuka (1999) perform calculations for the early collapse (before dissociation of H$_2$ in the core) of a 1\\,M$_\\odot$ star.  The luminosity stays below 0.1\\,L$_\\odot$, there is a core of 5\\,AU and the density goes like $r^{-2}$ in the envelope.  Their models should be applicable to HH108MMS and Q1 of the present paper which also have no mid--IR embedded source.  \\noindent In a few cases, the predicted theoretical density profiles could be checked by 1.3mm observations of the dust emission (Ward--Thompson et al.~1994, Andr\\'e et al.~1996, Ward--Thompson et al.~1999) and were generally corroborated.  In this method, the conversion of fluxes into absolute column densities depends on the dust absorption coefficient at 1.3\\,mm and the grain temperature, whereas the power law exponent of the density distribution is sensitive to the temperature gradient.  \\noindent On the other hand, our derivation of the optical depth and the ensuing density profile from absorption measurements is temperature independent.  ISOCAM observations similar to ours were also carried out by Bacmann et al.~(2000).  For the sources which they present, the maximum optical depth at the cloud center is much smaller than those for HH108MMS and Q1, so we probe deeper into the cloud, still the density profiles in the {\\it envelopes} roughly agree with what we derive.  In the {\\it cloud center} ($\\le 2000$AU), they find a flattening of the column density which we do not see.  Such a flattening is expected to occur in the central 1000AU of an isothermal sphere with temperature $T \\sim 10$K, radius of 10000AU and mass of 1\\Msun\\ simply by solving the hydrostatic equation (Bonnor 1956). \\noindent  \\noindent The measured mid infrared to millimeter dust extinction ratio of $ \\kappa_{14\\mu \\rm{m}} / \\kappa_{1300\\mu \\rm{m}} \\sim 470 $ should be compared with dust models of protostellar environments.  The ratio is a factor 4 lower than expected for dust in the diffuse medium and it indicates that the grains in the dense and cold environment of the protostar HH108MMS are, as expected (Ossenkopf 1993), of rather fluffy and composite nature.  The measured mid infrared to millimeter dust extinction ratio of HH108MMS is already tending towards somewhat elongated grain structures: In the fluffy composite but spherical dust model by Kr\\\"ugel \\& Siebenmorgen (1994) a ratio of $\\kappa_{14\\mu\\rm{m}} / \\kappa_{1300\\mu \\rm{m}} \\sim 1000$ is found. Because elongated particles are much better antennas at 1.3mm, they can give a lower value.  A proof that the dust in the absorbing cores and the IRAS source is indeed of elongated structure is found by the ISOCAM polarisation measurements.  \\noindent As the filament is seen in absorption against the background and assuming that the majority of the grains are still sub-micron sized particles, one may neglect dust scattered light at 14$\\mu$m.  Consequently the most plausible mechanism to produce the polarisation is dichroism.  For polarisation due to extinction of elongated spinning and aligned dust particles, the polarisation vectors (Figs.~\\ref{IRAS18331_pol_lw3.ps} and \\ref{IRAS18331_pol_lw10.ps}) indicate the magnetic field direction.  It appears that the magnetic field vectors are roughly aligned with the absorbing filament.  Such ordered fields are also detected from 850$\\mu {\\rm m}$ maps of thermal dust emission of prestellar cores (Ward-Thompson et al.~2000) and FIR polarimetry on the protostar IRAS20503+6006 (Clements et al.~1999)."
    },
    "0011/astro-ph0011583_arXiv.txt": {
        "abstract": "Gamma-ray bursts provide what is probably one of the messiest of all astrophysical data sets. Burst class properties are indistinct, as overlapping characteristics of individual bursts are convolved with effects of instrumental and sampling biases. Despite these complexities, data mining techniques have allowed new insights to be made about gamma-ray burst data. We demonstrate how data mining techniques have simultaneously allowed us to learn about gamma-ray burst detectors and data collection, cosmological effects in burst data, and properties of burst subclasses. We discuss the exciting future of this field, and the web-based tool we are developing (with support from the NASA AISR Program). We invite others to join us in AI-guided gamma-ray burst classification (http://grb.mnsu.edu/grb/). ",
        "introduction": "Understanding the physics of a class of astronomical objects depends on identifying intrinsic behaviors. When two or more subclasses are present, each subclass is defined in terms of its own intrinsic behaviors. The process of identifying behavioral characteristics is difficult when the objects' observed characteristics (or {\\em attributes}) overlap. Such is the case for cosmic gamma-ray bursts (GRBs), which have a large spread in observed attribute values. Some GRB attribute dispersion is intrinsic, some is caused by measurement error, some is due to systematic ({\\em e.g.} instrumental and sampling) biases, and some is caused by the presence of subclasses.  GRB subclass behaviors are difficult to delineate from other causes of attribute dispersion. Two GRB subclasses have been known to exist for some time \\cite{C74} \\cite{K93}, but it has been difficult to assign individual GRBs to a class because of attribute overlap. Class assignment has been complicated even more by the statistical clustering identification of a third GRB subclass \\cite{M98}; properties of this third subclass overlap those of the other two.  GRB classification can be aided by Knowledge Discovery in Databases (KDD) \\cite{D00}. The approach uses pattern recognition algorithms from the Artificial Intelligence (AI) branch of computer science to find behaviors indicative of subclasses. KDD offers a methodology by which meaningful information can be extracted from large volumes of data. The KDD process (Figure \\ref{fig:fig0}) is composed of data pre-processing and storage (data warehousing), data mining (clustering software), and scientific/logical assessment. Statistical and systematic effects (e.g. instrumentation and sampling biases) can be identified and even removed in the assessment step.  \\begin{figure}[htbp] \\begin{center}\\small \\unitlength=0.66mm \\linethickness{0.4pt} \\begin{picture}(245,80) \\put(5,20){\\dashbox{1}(45,50)[ct] {\\shortstack{\\\\ \\\\ Data Warehousing}}} \\put(12.5,30){\\framebox(30,30)[cc] {\\shortstack{Pre- \\\\ Processed \\\\ GRB Data}}} \\put(42.5,42){\\line(1,0){17}} \\put(42.5,48){\\line(1,0){17}} \\put(62.5,45){\\line(-1,+1){6}} \\put(62.5,45){\\line(-1,-1){6}} \\put(55,20){\\dashbox{1}(45,50)[ct] {\\shortstack{\\\\ \\\\ Data Mining}}} \\put(62.5,30){\\framebox(30,30)[cc] {\\shortstack{Classifier \\\\ Shell}}} \\put(93,42){\\line(1,0){8.5}} \\put(93,48){\\line(1,0){8.5}} \\put(104.5,45){\\line(-1,+1){6}} \\put(104.5,45){\\line(-1,-1){6}} \\put(105,30){\\framebox(30,30)[cc] {\\shortstack{Scientific \\\\ and \\\\ Logical \\\\ Assessment}}} \\put(135.5,42){\\line(1,0){5.5}} \\put(135.5,48){\\line(1,0){5.5}} \\put(144.5,45){\\line(-1,+1){6}} \\put(144.5,45){\\line(-1,-1){6}}  \\put(120,30){\\line(0,-1){15}} \\put(49.5,14){Iteration} \\put(120,15){\\line(-1,0){49}} \\put(47.5,15){\\line(-1,0){20}} \\put(27.5,15){\\vector(0,+1){15}}  \\put(145,30){\\framebox(30,30)[cc] {\\shortstack{GRB \\\\ Subclasses}}} \\put(0,10){\\dashbox{0.5}(140,70)[ct] {\\shortstack{\\\\ \\\\ Gamma-Ray Burst Classification Tool}}}  \\end{picture} \\end{center} \\vspace*{-10mm} \\caption{Gamma-Ray Burst Classification Process}\\label{fig:fig0} \\end{figure}  AI classifiers are typically {\\em supervised} or {\\em unsupervised}. Supervised classifiers require training {\\em instances} (data elements) in order to develop classification rules for unknown instances. Unsupervised classifiers try to subclassify a data set by searching for clusters in multidimensional attribute spaces.  We are developing a web-based tool \\cite{Hag00} for the classification of GRB data (http://grb.mnsu.edu/grb/). The tool contains a preprocessed GRB database, AI classifiers, and data visualization software. In this manuscript we describe some of our initial scientific results concerning GRB data mining with this tool. Additional results have been published elsewhere \\cite{H00}, \\cite{R00}, \\cite{Hak00}. ",
        "conclusions": "We have demonstrated that data mining techniques can aid the interpretation of scientific data, even with complex and ambiguous GRB data.  Data mining demonstrates that some Class 1 (Long) GRBs can develop Class 3 (Intermediate) characteristics via a combination of the hardness intensity relation and the fluence duration bias. Class 3 (Intermediate) GRBs do not appear to represent a separate source population, although they cluster in the duration, fluence, hardness, attribute space. Class 2 (Short) GRBs do appear to represent a separate source population."
    },
    "0011/astro-ph0011530_arXiv.txt": {
        "abstract": "We study spectroscopically determined iron abundances of $642$ solar-type stars to search for the signature of accreted iron-rich material. We find that the metallicity [Fe/H] of a subset of $466$ main sequence stars, when plotted as a function of stellar mass, mimics the pattern seen in lithium abundances in open clusters. Using Monte Carlo models we find that, on average, these stars have accreted $\\sim0.4M_\\oplus$ of iron while on the main sequence. A much smaller sample of $19$ stars in the Hertzsprung gap, which are slightly evolved and whose convection zones are significantly more massive, have lower average [Fe/H], and their metallicity shows no clear variation with stellar mass. These findings suggest that terrestrial-type material is common around solar type stars. ",
        "introduction": "The discovery that at least $6-8\\%$ of solar type stars harbor Jupiter-mass or larger bodies, often in small, eccentric orbits \\citep{mq, marcy, butler, mcm} shows that planets are not exceptionally rare. The transiting planet orbiting HD 209458 has a mass of $0.69$ Jupiter masses, and a radius about 1.4 times that of Jupiter. Clearly it is a gas giant like Jupiter or Saturn; the minimum masses of the other known objects suggest that they are also gas giants. The Doppler technique used for most of the discoveries cannot find terrestrial mass objects in AU scale orbits. What fraction of solar type stars have terrestrial-type planets?  How could an observer with our current technology located on a star in the solar neighborhood decide if there were any terrestrial-type bodies orbiting the sun?  More generally, how could such an observer estimate the fraction of solar type stars having companions of terrestrial (as opposed to gas giant) compositions? In this paper we explore one possible way of addressing the latter question, stellar-mass dependent photospheric metallicity trends.  In section \\ref{pollution} we argue that a few Earth masses of rocky/icy material have accreted onto the Sun over its lifetime. In section \\ref{evidence} we discuss one possible observational tracer of similar accretion occurring on other solar type stars. In section \\ref{observations} we examine a sample of stars with spectroscopically measured photospheric iron abundances \\be % {\\rm [Fe/H]}\\equiv\\log\\left[{f_{Fe}\\over f_{Fe,\\odot}}\\right], \\ee % where $f_{Fe}$ is the mass abundance of iron in the photosphere of the star, and $f_{Fe,\\odot}\\approx 1.3\\times10^{-3}$ is the the mass abundance of iron in the photosphere of the sun.  We compare the observations to Monte Carlo models of stellar pollution in section \\ref{gamble}. We discuss the implications of our results in section \\ref{discussion} and present our conclusion in the final section. ",
        "conclusions": "We analyzed some 642 stars from the Cayrel de Strobel catalogue having spectroscopically determined metallicities, and well determined HIPPARCOS parallaxes. Using a large grid of stellar models, we determined the age and mass of the stars in our sample. We then examined the variation of metallicity with stellar age and mass.  We have found striking mass-dependent variations in photospheric iron abundances of main sequence solar mass ($0.8-1.8M_\\odot$) stars in the solar neighborhood. These variations mimic the variations seen in lithium abundances. With somewhat less confidence, because of the small sample size, we find that the iron abundances of Hertzsprung-gap stars are on average lower than those of main sequence stars, and that the metallicities of these slightly evolved stars have no mass dependence. These results are consistent with the accretion of an average of $\\sim0.4M_\\oplus$ of iron onto the surface of main sequence stars. This strongly suggests that terrestrial-type material is common around solar type stars in the solar neighborhood."
    },
    "0011/astro-ph0011195_arXiv.txt": {
        "abstract": "Observations of the cosmic microwave sky are revealing the primordial non-uniformities from which all structure in the Universe grew. The only known physical mechanism for generating the inhomogeneities we see involves the amplification of quantum fluctuations during a period of inflation. Developing the theory further will require progress in quantum cosmology,  connecting inflation to a theory of the initial state of the Universe. I discuss recent work within the framework of the Euclidean no boundary proposal, specifically classical instanton solutions and the computation of fluctuations around them. Within this framework, and for a generic inflationary theory,  it appears that an additional anthropic constraint is required to explain the observed Universe. I outline an attempt to impose such a constraint in a precise mathematical manner. ",
        "introduction": "The maps of the cosmic microwave sky provided by recent experiments  \\cite{1} signal an important breakthrough  in cosmology. The structure revealed provides our most powerful probe of the structure in the Universe at early times, and of its current geometry. The data is consistent the Universe  being nearly spatially flat, with a scale invariant spectrum of Gaussian-distributed irregularities in which the overall density fluctuates but the equation of state does not. This simplest form was anticipated on phenomenological grounds as far back as the sixties \\cite{2}, but later emerged as the prediction of simple models of inflation \\cite{3,4,5,6}. The detailed agreement between theory and observations lends impressive support to the view that the Universe is, in some respects at least, remarkably simple.  In inflationary theories, inhomogeneities of the required form arise a side-effect of the exponential expansion invoked to explain the size and flatness of the Universe. Quantum mechanical vacuum fluctuations in the inflaton field are stretched to large scales during inflation \\cite{5,6}, later to re-enter the Hubble radius and seed large scale structure. This magnificent mechanism links quantum mechanics and gravity to the observed structure of the Universe. It provides a direct observational probe of quantum gravity, albeit at leading (Gaussian) order.   The underlying theory is however only provisional, because we still have no compelling model of inflation and because a complete theory of quantum gravity is lacking. Nevertheless the success of the inflationary explanation of the structure in the cosmic microwave sky forces us to take very seriously the idea that quantum mechanics governed the formation of the Universe we see. This moves the topic of quantum cosmology to centre stage. It is a field that {\\it must} be developed if we are attain a deeper level of understanding in cosmology.  The prevalent view is that cosmology will over the next period be observation-driven. It is true that dramatic observational developments are certain, and there will be a lot of phenomenology to do. However if we mean to achieve more than simply measuring the properties of the universe, theoretical developments are also crucial. The great freedom there is in current approaches to phenomenological models of inflation and of the early Universe, combined with the relatively small number of observations (even including the cosmic microwave sky) leads me to conclude that if we are to get anywhere, mathematical requirements of completeness and consistency must play an increasing role. Their power is already convincingly demonstrated by the fact that not one of the existing inflationary models has been made sense of beyond leading (Gaussian) order.   In this talk I will concentrate on one particular incompleteness of inflationary theory. The question is why inflation started in the first place. Vilenkin, Linde and Guth argue that the need for a theory of initial conditions is side-stepped because inflation is self-sustaining to the future (i.e. `semi-eternal')\\cite{eternal}. Whilst accepting that some theory of the initial conditions is needed at some level, the claim is that the details of that theory are actually irrelevant, since there would typically be an infinite amount of inflation between ourselves and the beginning, during which details of the initial conditions would have been erased. In this view discussions of the initial conditions prior to inflation are fairly unimportant, since `almost anything' will do. In this talk I criticise the calculational methods which have been used to reach this conclusion, and explain a different point of view according to which inflation is {\\it a priori} very unlikely and in any case only a brief episode in our past.  One can easily imagine an infinite number of possible ansatzes for the initial conditions of the Universe. At the present stage of development of theory, many of these might be perfectly consistent with observation. But among the existing proposals, I think the Euclidean no boundary proposal \\cite{hh} is appealing because it is based on simple and general ideas which have a rationale beyond cosmology. In a strong sense I think it is `the most conservative thing you can do'. Of course it may well fail precisely because it is too conservative. Space and time may be emergent rather than fundamental properties. Describing the Universe as a manifold may not be appropriate to its early moments. Nevertheless precisely because the no boundary proposal is {\\it not} just cooked up to make inflation work, its failings and limitations may teach us something deeper about what is in fact required.   In this lecture I focus on one particular approach to the no boundary proposal, using a set of classical instanton solutions of the `no boundary' form \\cite{ht,th}. Our work has focussed on using these to compute the the complete Gaussian correlators \\cite{gt,het,hht,ght}, in a generic inflationary model (for related work see \\cite{otha,othb}). These generic instantons are unusual in several respects. In the original `Einstein frame' they exhibit a curvature singularity which in the Lorentzian spacetime is `naked'. Nevertheless the Euclidean action for the solutions is finite, suggesting they should contribute to the Euclidean path integral. Second, there is a one parameter family of solutions, with differing action. This is inconsistent with their being true stationary points of the action, but they may nevertheless be legitimate as `constrained instantons' which are an established tool in other contexts \\cite{affleck}. A second feature is that the quantum fluctuations about these solutions are well defined to leading order, since the Euclidean action selects a unique (Dirichlet) boundary condition at the singularity. This allows one to make unambiguous predictions for the microwave anisotropies in any one of these solutions, which Gratton, Hertog and I have recently computed \\cite{gt,het,ght}. (See also \\cite{hht}).  I then describe recent work by Kirklin, Wiseman and myself \\cite{ktw} showing how the singularity apparent in the original `Einstein frame' may actually be removed by a change of field variables. Thus the singularity is really only a coordinate singularity on superspace. One still has the problem that the scalar field potential energy diverges at the singularity, and the scalar field equation remains ill defined there. That problem is overcome if one defines the theory as the limit of a family of theories in which $V(\\phi)$ takes a particular form as $\\phi$ tends to infinity, but approximates the actual potential at arbitrarily large values of $\\phi$. The limit is insensitive to the precise form of the regularised potential at large $\\phi$.  The regularised construction allows one to give an improved formulation of the constraint at the singularity, which allows for a detailed computation of the spectrum of homogeneous modes \\cite{gtnew}. The latter, we argue, are actually crucial to a proper interpretion of the instantons.  In the regularised theory all quantities appearing are finite. The classical field equations are also satisfied everywhere if one views the instantons as topologically $RP^4$ rather than $S^4$. The same construction is applied to instantons possessing two singularities in the Einstein frame. When the singularities are `blown up' one obtains a regular manifold which is a four dimensional analogue of the Klein bottle.  Finally I discuss the problem of negative modes, complicated  by the `conformal factor' problem of Euclidean general relativity, due to the non-positive character of the Euclidean action. In the regular variables (and with the regularised potential), the problem of negative modes about singular instantons is well defined, and I briefly mention some of our latest results \\cite{gtnew} indicating that the most interesting singular instantons (describing Universes with a large amount of inflation) do not possess any negative modes, in contrast to the non-singular instantons of Coleman and De Luccia, and Hawking and Moss. I comment on how this observation might help to resolve the `empty Universe problem', if one performs a projection onto Universes containing the observer. ",
        "conclusions": ""
    },
    "0011/astro-ph0011085_arXiv.txt": {
        "abstract": " ",
        "introduction": "Surface photometry is one of the oldest techniques in modern Astronomy. The first attempt of surface photometry was done in Helwan observatory in Egypt by Reynolds 1913. In quanititative analysis of the structure of galaxies, some of the two-dimensional data are reduced to one-dimensional luminosity profiles. The shape of the surface brightness profile took more attention and has been put forward as a new extragalactic distance indicator (c.f. Binggeli \\& Jerjen 1998).  NGC 3077 is a member galaxy  of the M81 interacting group of galaxies. It lies  46$^{\\prime}$ south-east of the spiral galaxy M81. Two HI bridges connect it with its neighbouring galaxies M81 and M82 (cf. Yun et al. 1994, Thomasson and Donner 1993). It was classified as an Irr II galaxy (Sandage 1961). The centre of NGC 3077 is dusty and exhibits many compact blue knots, in addition to condensed HII regions embedded in diffuse ionized gas (Barbieri et al. 1974). In this paper we performed detailed continuum surface photometry to investigate the morphology and the colour distribution of the galaxy NGC 3077. ",
        "conclusions": "A study of the brightness and colour distribution of NGC 3077 shows that this galaxy consists of a blue disturbed centre embedded in a smooth elliptical structure; this structural type is common in  irregular galaxies. The lack of symmetry in the inner region is due to the presence of a young stellar population reddened by dust.  The colours of this young  population indicate ages between 4 and 100 Myr, whereas the mean age of the dominant population in the outer region of NGC 3077 is 3$\\cdot 10^9$ yr. The age range of the young central population suggests that the ongoing encounter with M81, which commenced a few hundred million years from now (eg. Thomasson and Donner 1993), serves as a trigger for the star formation in the galaxy's core."
    },
    "0011/astro-ph0011566_arXiv.txt": {
        "abstract": "We report optical, radio and X--ray observations of a new distant blazar, PMN~J0525-3343, at a redshift of 4.4. The X-ray spectrum measured from ASCA and $Beppo$SAX flattens below a few keV, in a manner similar to the spectra of two other $z>4$ blazars, GB~1428+4217 ($z$=4.72) reported by Boller et al and RXJ~1028.6-0844 ($z$=4.28) by Yuan et al. The spectrum is well fitted by a power-law continuum which is either absorbed or breaks at a few keV. An intrinsic column density corresponding to $2\\times 10^{23}$H-atoms\\ \\psqcm at solar abundance is required by the absorption model. This is however a million times greater than the neutral hydrogen, or dust, column density implied by the optical spectrum, which covers the rest-frame UV emission of the blazar nucleus.  We discuss the problems raised and suggest that, unless there is intrinsic flattening in the spectral distribution of the particles/seed photons producing X--rays via inverse Compton scattering, the most plausible solution is a warm absorber close to the active nucleus. ",
        "introduction": "Active Galactic Nuclei (AGN) at high redshift are powerful tools with which to study the evolution of massive black holes and of their young galaxy hosts. A particularly interesting class is formed by several high redshift ($z>4$), X-ray bright, radio--loud quasars (Fabian et al 1997, 1998; Moran \\& Helfand 1997; Zickgraf et al 1997; Hook \\& McMahon 1998) which present characteristics typical of blazars. One such source, GB~1428+4217, is variable in both the X-ray and radio bands and has a spectral energy distribution which peaks at hard X-ray energies (Fabian et al 1998).  Recently this object (Boller et al 2000) and RXJ1028.6-0844 (Yuan et al 2000) have been found to show spectral flattening which is interpreted as due to X-ray absorption, implying an absorbing column density of $\\sim 1.5\\times10^{22}$\\psqcm and $\\sim 2.1\\times 10^{23}$\\psqcm, respectively, if intrinsic. We report here on the discovery and study of a similar object, PMN~J0525-3443 at $z=4.4$, which shows similar spectral flattening.  Fiore et al (1998) have earlier found from ROSAT data an apparent systematic decrease with redshift of the spectral slope of the soft X--ray emission of radio-loud quasars from local to $z\\sim 3.9$ objects.  The change in spectra with redshift might be associated with an increase with redshift in the amount of absorbing (intrinsic or external) gas. Cappi et al (1997) find evidence for absorption using ASCA data for a small sample of high-redshift ($1.2<z<3.4$) radio-loud quasars. Confirmation of a systematic spectral flattening has been also reported by Reeves \\& Turner (2000) for a significant number of radio--loud sources observed by ASCA up to $z$=4.2. Understanding the nature and evolution of this spectral change will be of great importance for our knowledge of gaseous environments of powerful quasars, their intrinsic properties as well as their cosmological evolution and the radio--loud quasar/galaxy connection.  The object under study here, PMN J0525-3343, was identified by McMahon and Hook as part of their program to identify high redshift radio--loud quasars (Hook \\& McMahon 1998).  The quasar was selected as a flat spectrum (i.e. $\\alpha < $ 0.5; $F_{\\nu} \\propto \\nu^{-\\alpha}$) radio source using the 4850 MHz Parkes-MIT-NRAO (PMN) survey (Wright et al 1996) and 1.4 GHz NRAO VLA Sky Survey (NVSS, Condon et al 1998).  The 5.0GHz and 20cm fluxes are 210$\\pm$18 and 188.7$\\pm$5.7 mJy, respectively, giving a spectral index $\\alpha_{1.49-5.0 GHz} = -0.1$.  The radio sources were matched against the APM UKST catalogue using the NVSS radio position and it was identified as a red stellar object with R=18.5 and colour B--R=2.80. Subsequent optical spectroscopy identified the radio source as a quasar with z=4.4.  It was confirmed to be an X--ray source with the ROSAT HRI and we have since observed it with ASCA and $Beppo$SAX. We present here X-ray fluxes and spectra, together with measurements of its radio and optical spectra.  PMN~J0525-3443 is the second most distant object in the X--ray bright, radio--loud class of quasars.  The optical, radio and X--ray data analysis procedure and results are reported in Sections~2, 3 and 4, respectively, while discussion and conclusions are presented in Sections~5 and 6.  $H_0$=50 km s$^{-1}$ Mpc$^{-1}$ and $q_0$=0.5 have been used throughout the paper. ",
        "conclusions": "We have studied the X-ray spectral properties of the high redshift radio--loud quasar PMN~J0525-3343, by means of ROSAT, ASCA and $Beppo$SAX observations. We have found that the spectrum flattens at the lowest energies, indicating either a spectral flattening or a significant absorption in excess of the Galactic one.  This finding is consistent with and extends the range in redshift of the systematic trend of increasing flattening with increasing redshift found previously for radio--loud sources. Several possibilities have been discussed.  The most plausible hypotheses include either an intrinsic flattening of the soft X--ray spectrum, originating from the shape of the energy distribution of particles and/or seed photons involved in the production of the Compton component, or the presence of warm absorbing gas in the nuclear region. If this is similar to the absorber found in local (Seyfert) objects, its distance from the central source would be of order $\\sim 10^{16}$ cm. Evidence for CIV absorption in the spectrum supports the warm absorber hypothesis. The equivalent width of the absorption and value of $\\alpha_{\\rm ox}$ for PMN J0525-3343 fall on the correlation of Brandt et al (2000; Figure 4). Indeed the objects neighbouring PMN J0525-3343 in that correlation are observed to have X-ray warm absorbers. The velocity difference between the CIV absorption and emission are unlikely, in an object at $z=4.4,$ to be due to an intervening galaxy in a massive cluster of galaxies, since such clusters would not have then formed. Possibly the absorption line redshift is that of the quasar nucleus, and the broad emission line profiles have been altered by further, unresolved, absorption into peaking at a lower redshift.  In the former case, the column-redshift dependence could be accounted for by the decreasing contribution of synchrotron self--Compton emission at the softest energies due to an increasingly intense external photon field. In the latter, one might envisage an increase in the amount of nuclear gas and/or of ionizing radiation with $z$. Interestingly, it has been suggested (e.g. Reeves \\& Turner 2000 and references therein) from the observed characteristics of the spectra of radio--quiet sources, that the key parameter regulating their spectral properties is the mass accretion rate in Eddington units: high luminosity sources would thus be characterized by an increasing ionization state of the reflecting disc and possibly of any nuclear absorber. A similar trend might be at the origin of the increasing photon density required by the above scenario in radio--loud objects.  Future observations with Chandra and XMM-Newton, and in the optical, will enable us to discriminate between the models for these sources.  Clearly, this issue needs more investigation, since it represents a potential means for understanding the conditions and evolution of the nuclear/galactic environment and the galaxy--quasar connection.  In particular future X--ray observations with wide energy coverage and good spectral resolution as provided by a significant number of high redshift radio--loud objects, will allow us to a) establish the origin of the spectral flattening, determining the abundance and ionization state of any absorber and -- if variability or clear features are detected -- its spatial location and redshift/velocity and b) assess its evolutionary properties. Information from the optical--UV spectrum are also of prime importance, as the amount of (or limit on) the extinction can be directly compared to the X--ray one and any line absorption feature would determine the absorbing gas velocity, although photometry alone might not be sufficient to distinguish between the steep non--thermal continuum typical of blazars and a reddened quasar spectrum."
    },
    "0011/astro-ph0011420_arXiv.txt": {
        "abstract": "We study the relation between gaseous absorbing column density (\\nha), infrared colors and detectability of the broad lines in a large sample of Seyfert 2 galaxies(Sy2s). We confirm that Sy2s without polarized broad lines tend to have cooler 60$\\mu$m/25$\\mu$m colors; this correlation was previously ascribed to the effect of obscuration towards the nuclear region. We find some evidence that Sy2s without polarized broad lines have larger absorbing column density (\\nha) and that a fraction of them are characterized by dust lanes crossing their nuclei. However, we find that the IR colors do not correlate with \\nha, in disagreement with the obscuration scenario. Also, Sy2s without polarized broad lines follow the same radio-FIR relation as normal and starburst galaxies, at variance with Sy2s with polarized broad lines. These results indicate that the lack of broad lines in the polarized spectrum of Sy2s is mostly due to the contribution/dilution from the host galaxy or from a circumnuclear starburst, though at a lower extent the obscuration toward the nuclear region also plays a role. ",
        "introduction": "According to the standard unification model, Seyfert 1 and 2 galaxies (Sy1s and Sy2s hereafter) are intrinsically the same objects and the absence of broad lines in Sy2s is ascribed to the obscuration by a pc-scale dusty torus oriented along the line of sight (see the reviews by  Antonucci \\cite{antonucci_rev} and V\\'eron-Cetty \\& V\\'eron  \\cite{veron}). The observational evidence for this model includes the detection of polarized broad emission lines in some Seyfert 2 galaxies (Antonucci \\& Miller \\cite{antonucci}; Tran \\cite{tran}; Moran et al. \\cite{moran}), the detection of broad lines in the infrared spectrum of some Sy2s (Ruiz et al. \\cite{ruiz}; Veilleux et al. \\cite{veilleux}; Rix et al.  \\cite{rix}) and the detection of a prominent photoelectric cutoff in the X-ray spectra of Sy2s indicating the presence of large columns of gas along the line of sight (Koyama et al. \\cite{koyama}; Awaki et al. \\cite{awaki}; Maiolino et al. \\cite{maiolino98}; Risaliti et al. \\cite{risaliti}).  The spectropolarimetric observations of different samples of Seyfert 2 galaxies indicate that {\\it only} about 40 \\% of Sy2s show broad lines in their polarized spectra (eg. Heisler et al. \\cite{heisler}), although such surveys are probably biased since pre-selection was done according to the broad-band polarization.  According to the suggestion of Heisler et al. (\\cite{heisler}), the detectability of a hidden BLR through spetropolarimetry in Sy2s is related to the inclination of the torus which, in turn, is related to the 60 $\\mu$m to 25 $\\mu$m flux ratio, $s_{60\\mu m}/s_{25\\mu m}$. More specifically, in those Sy2s showing polarized broad lines (PBL) the torus should be oriented more face-on so that the scattering medium is less obscured by the torus itself. This model would also explain the correlation between IR colors ($s_{60\\mu m}/s_{25\\mu m}$) and detectability of the PBL. In particular, when the torus is observed close to pole-on, the PBL should be more easy to detect and the IR color should be hotter since we are observing the hotter dust emitting region, in agreement with what is observed.  However, more recently Alexander (\\cite{alexander}) compared the absorbing column densities inferred from the hard X-rays with the detectability of the PBL and found no correlation. This result is in contrast with Heisler's et al. (\\cite{heisler}) model, which would predict a higher absorbing column for Sy2s without PBL. Alexander (\\cite{alexander}) suggests that the relation between detectability of the PBL and IR colors is indirect: the contribution from the host galaxy would both make the IR color cooler and would also dilute the nuclear optical spectrum making more difficult the detection of scattered polarized light.  The simple model of the obscuring torus has been subject to various modifications. In particular, while the pc-scale torus is probably responsible for the huge absorbing columns (\\nh $> 10^{24}$ cm$^{-2}$) observed in several Sy2s, observational evidence was also found for a larger scale ($\\sim$ 100 pc) obscuring medium with lower absorbing column density (\\nh $\\sim$ a few times $10^{22}$ cm$^{-2}$, Granato et al. \\cite{granato}; Matt \\cite{matt}; Maiolino \\cite{maiolino} and references therein).  In this paper we expand the work done by Alexander (\\cite{alexander}) by enlarging the sample of Sy2s for which information on both the absorbing \\nh and on the detection of PBL is available. We also seek additional constraints on the nature of the circumnuclear scattering and absorbing medium by comparing the mid- and far-IR colors with \\nh and the detectability of PBL with the nuclear morphology and with the radio power.  The paper is organized as follows. In Section 2 we present our sample of Seyfert 2 galaxies with spectropolarimetric observations, the results are given in Section 3. We discuss our results and their implications in Section 4 and summarize our conclusions in Section 5. ",
        "conclusions": "In this paper, we collect 40 Seyfert 2 galaxies having both spectropolarimetric observations and a measure of the absorbing \\nh obtained by means of their X-ray spectrum. Out of these 40 objects 22 show broad lines in polarization (most likely ascribed to scattering of the broad line region) and 18 do not. For these objects we also analyzed the relation of \\nh and detectability of the broad lines with the mid- and far-IR colors. We confirm previous claims that Sy2s without polarized broad lines have cooler IR colors. We studied additional diagrams to test two scenarios proposed to explain such a correlation and that, more specifically, ascribe this effect either to obscuration of the nuclear region or to the contribution/dilution from circumnuclear starburst activity and from the host galaxy.  We found a marginal evidence for Sy2s showing broad lines in polarization to have lower \\nh. We also find that about 30\\% of the Sy2s without polarized broad lines show dust lanes crossing the nucleus, while none of the Sy2s with polarized broad lines show evidence for such nuclear dusty structures. These results suggest that the obscuration towards the nuclear regions (hence towards the scattering mirror) plays a role in hiding the polarized broad lines, at least in some objects.  On the other hand, we find that the absorbing column density does not correlate with the IR colors and, in particular, these quantities do not follow the relation expected in the case of absorption of the IR emitting region by the dust associated to the observed \\nha, therefore indicating that the distribution of IR colors is not to be ascribed to obscuration effects. Also, we find that Sy2s without detection of polarized broad lines follow the same radio--FIR relation as normal and starburst galaxies, at variance with Sy2s with polarized broad lines which tend to spread towards higher radio luminosities. These findings support previous claims that the relation between far-IR colors and detectability of the polarized broad lines in Sy2s is mostly related to dilution of the IR and optical light by a circumnuclear starburst  or by the host galaxy.  Summarizing, our results indicate that the lack of broad lines in the polarized spectrum of Sy2s is mostly due to the contribution/dilution from the stellar component, though at a lower extent the obscuration towards the nuclear region also plays a role."
    },
    "0011/astro-ph0011346_arXiv.txt": {
        "abstract": "We report the results of a survey of the Small Magellanic Cloud (SMC) for radio pulsars conducted with the 20-cm multibeam receiver of the Parkes 64-meter telescope. This survey targeted a more complete region of the SMC than a previous pulsar search and had an improvement in sensitivity by a factor of about two for most pulsar periods. This survey is much more sensitive to fast young pulsars (with $P \\approxlt 100$ ms) and is the first survey of the SMC with any sensitivity to millisecond pulsars. Two new pulsars were discovered in the survey, one of which is located within the SMC. The number of pulsars found in the survey is consistent with the expected number derived using several methods.  We also report the serendipitous discovery of a new pulsar in the 30 Doradus region of the Large Magellanic Cloud (LMC). These discoveries bring the total number of rotation-powered pulsars currently known in the Magellanic Clouds to eight. We have also made refined timing measurements for the new discoveries as well as for three previously known LMC pulsars.  The age distribution of luminous Magellanic Cloud pulsars supports the conjecture that pulsars younger than about 5 Myr are more luminous on average than older pulsars. ",
        "introduction": "The Magellanic Clouds contain the most distant population of radio pulsars observable with current technology. Since the distances to the Magellanic Clouds are large, only the most luminous pulsars are detectable. However, despite the small number of pulsars currently known, these pulsars are among the most interesting in the pulsar population. Of the six pulsars discovered previously, one is in an unusual binary system and two are younger than 10 kyr. For comparison, of over 700 Galactic pulsars cataloged prior to the commencement of the Parkes Multibeam Pulsar Survey (see below), only one is in a similar binary system, and only five Galactic pulsars are younger than 10 kyr.  Surveying the Magellanic Clouds for radio pulsars is difficult due to the large surface area to be covered and their large distances. Long integrations using a telescope with a large collecting area are needed.  With its southern latitude, the Parkes 64-meter radio telescope in NSW, Australia is the most appropriate telescope for conducting such searches.  The difficulty of obtaining sufficient sensitivity has been somewhat alleviated by the installation of the 20-cm multibeam receiver at Parkes (\\cite{swb+96}). This receiver is simultaneously sensitive to 13 distinct locations on the sky which can be interleaved in separate pointings to give complete and efficient spatial coverage of large regions.  The Parkes Multibeam Pulsar Survey is currently using this receiver to search for pulsars in the Galactic plane, and that survey has discovered 580 pulsars to date (\\cite{cam00}; \\cite{man00}; \\cite{lcm+00}). This demonstrates that the multibeam system is capable of finding faint pulsars and is an appropriate instrument to search for pulsars in the Magellanic Clouds.  Only one full-scale pulsar survey of the Magellanic Clouds has been conducted previously. McConnell et al.~(1991)\\nocite{mmh+91} searched for pulsars using Parkes at a wavelength of 50 cm.  This survey effort was split into three phases with different observing parameters used for each phase of the survey. Although the region surveyed in each phase was not described, the third phase of the survey was used for five of the seven years of the effort, and we use this phase for a sensitivity comparison with our survey. In this phase, they used a center frequency of 610 MHz with a total bandwidth of 60 MHz split into 24 2.5-MHz contiguous frequency channels. Two orthogonal polarizations were used with a sampling time of 5 ms and a 5000 s dwell time for each pointing. The second phase of their survey had a comparable sensitivity to the third phase, but the first phase had significantly reduced sensitivity to fast pulsars due to the much larger sampling time of 40 ms used.  In their data reduction procedure, McConnell et al.~(1991)\\nocite{mmh+91} assumed a dispersion measure (DM) of 100 pc cm$^{-3}$ in an initial sub-band dedispersion in their analysis. Magellanic Cloud pulsars found to date have DMs which are roughly centered on this value, but the observed DM range is from about 70 pc cm$^{-3}$ to 150 pc cm$^{-3}$. Thus, their sensitivity to pulsars with DMs far from 100 pc cm$^{-3}$ was significantly reduced. McConnell et al.~(1991)\\nocite{mmh+91} also targeted the optical bar and the southern and western parts of the Small Magellanic Cloud (SMC) (see their Figure 1). Our survey had more complete coverage in the northern and eastern regions of the SMC.  McConnell et al.~(1991)\\nocite{mmh+91} reported the discovery of four pulsars, three in the Large Magellanic Cloud (LMC) and one in the SMC.  We report the results of a more sensitive and complete survey for pulsars in the SMC in which we have discovered one new SMC pulsar and one foreground pulsar. Below we describe the observations and data reduction procedure for our survey and the subsequent timing observations. We also report on the serendipitous discovery of a pulsar in the 30 Doradus region of the LMC, as well as on timing observations of three previously known LMC pulsars.  We discuss our results and summarize the currently known Magellanic Cloud pulsar population. The two new Magellanic Cloud pulsars found in our survey, combined with the four pulsars found by McConnell et al. (1991)\\nocite{mmh+91} and the two known X-ray pulsars in the LMC, bring the total number of known rotation-powered pulsars in the Magellanic Clouds to eight. ",
        "conclusions": "We have conducted a survey of the SMC for radio pulsars and have discovered two new pulsars, one of which is located within the SMC. We present timing results for both of these pulsars as well as refined values for three previously known LMC pulsars. We have also discovered serendipitously a new pulsar in the LMC, bringing the total number of known rotation-powered pulsars in the Magellanic Clouds to eight. The results of our survey are consistent with the expected number of detectable SMC pulsars estimated using several methods. However, the small number of pulsars currently known in the SMC prevents significant constraints on the parameter assumptions made in these estimates. The luminosity distribution of the Magellanic Cloud pulsar population on the $P$-$\\dot{P}$ diagram is consistent with the Galactic distribution of luminous pulsars and supports the conclusion of Taylor, Manchester, \\& Lyne~(1993)\\nocite{tml93} that pulsars younger than $\\sim$ 5 Myr are more luminous on average than older pulsars. The DM distribution of the newly discovered pulsars is consistent with the known pulsar population in the Magellanic Clouds. Both of the known SMC pulsars are in a narrow DM range, and if additional SMC pulsars are found in the future which are also confined to this narrow range, this will be at odds with the SMC depth estimated from Cepheid observations. Significantly increasing the sensitivity of searches for pulsars in the Magellanic Clouds is not practical with current technology given constraints on telescope time, etc. However, next-generation instruments such as a square-kilometer array should greatly increase the number of known pulsars in the Magellanic Clouds and will enhance and refine the conclusions drawn in this paper.  \\nocite{cra00}"
    },
    "0011/astro-ph0011493_arXiv.txt": {
        "abstract": "We have used a suite of simplified spectrophotometric spiral galaxy evolution models to argue that there are substantial variations in stellar mass-to-light ratios (\\mlsn) within and among galaxies, amounting to factors of between 3 and 7 in the optical, and factors of 2 in the near-infrared. Our models show a strong correlation between stellar \\ml and the optical colors of the integrated stellar populations. Under the assumption of a universal spiral galaxy IMF, relative trends in model stellar \\ml with color are robust to uncertainties in stellar population and galaxy evolution modeling, including the effects of modest bursts of star formation. Errors in the dust reddening estimates do not strongly affect the final derived stellar masses of a stellar population. We examine the observed maximum disk stellar \\mls of a sample of spiral galaxies with accurate rotation curves and optical and near-infrared luminosity profiles. From these observed maximum disk \\mls we conclude that a Salpeter Initial Mass Function (IMF) has too many low-mass stars per unit luminosity, but that an IMF similar to the Salpeter IMF at the high-mass end with less low-mass stars (giving stellar \\mls 30\\% lower than the Salpeter value) is consistent with the maximum disk constraints.  Trends in observed maximum disk stellar \\mls with color provide a good match to the predicted model relation, suggesting that the spiral galaxy stellar IMF is universal and that a fraction of (particularly high surface brightness) spiral galaxies may be close to maximum disk. We apply the model trends in stellar \\ml with color to the Tully-Fisher (TF) relation. We find that the stellar mass TF relation is relatively steep and has modest scatter, and is independent of the passband and color used to derive the stellar masses, again lending support for a universal IMF. The difference in slope between the optical (especially blue) and near-infrared TF relations is due to the combined effects of dust attenuation and stellar \\ml variations with galaxy mass. Assuming the HST Key Project distance to the Ursa Major Cluster and neglecting the (uncertain) molecular gas fraction, we find that the baryonic TF relation takes the form $M_{\\rm baryon} \\propto V^{3.5}$ (with random and systematic 1$\\sigma$ slope errors of $\\sim$ 0.2 each) when using a bisector fit and rotation velocities derived from the flat part of the rotation curve. Since we have normalized the stellar \\mls to be as high as can possibly be allowed by maximum disk constraints, the slope of the baryonic TF relation will be somewhat shallower than 3.5 if all disks are substantially sub-maximal. ",
        "introduction": "The stellar mass-to-light ratio (\\mln) is an important parameter in astrophysics as it allows translation between photometry and dynamics.  The stellar \\ml has a direct bearing on two hotly debated areas in spiral galaxy research: the appropriate stellar \\mls to be used for spiral galaxy rotation curve decompositions, and the passband-dependent slope of the galaxy magnitude-rotation velocity relation \\citep[TF relation hereafter]{tully77}. In this paper, we address the stellar \\mls of spiral galaxies, briefly explore the implications of our results for rotation curve decompositions and investigate in more depth the slope of the TF relation.  There is presently much interest in decomposing spiral galaxy rotation curves into contributions from the gaseous, stellar and dark matter contents \\citep[e.g.][]{verheijen97,deblok98}. The primary motivation for this interest is that, in principle, the structure of dark matter halos can be determined from spiral galaxy rotation curves {\\it if} the contribution from gas and stars can be properly understood. In turn, the structure of dark matter halos is a strong constraint on dark matter halo formation models \\citep[e.g.][]{moore98,navarro00a}. The main challenge in determining the dark matter contribution to a given rotation curve is our ignorance regarding plausible values of the stellar \\mln: the gas contribution is typically well-understood and relatively small \\citep{verheijen97,swaters00}. The situation is degenerate enough that many rotation curves can be equally well-fit by models in which the central parts of the rotation curve are dominated entirely by stellar mass or by dark matter \\citep[e.g.][]{vanalbada85,swaters99}.  In order to resolve this degeneracy, some independent constraints on stellar \\mlsn, and their variations with radius and galaxy properties, are required.  The implications of the stellar \\ml for the TF relation are no less important.  The TF relation relates the integrated luminosity in a given passband to the global dynamics of the galaxy and its dark matter halo. The dust-corrected TF relation has a slope which steepens towards redder passbands \\citep[going between $L \\propto V^3$ or shallower in the optical to $L \\propto V^4$ in the near-infrared; ][]{verheijen97,tully98} indicating that there is a trend in color and stellar \\ml with galaxy mass.  This change in slope with passband can considerably weaken the power of the TF relation as a test of galaxy formation and evolution models \\citep[such as those by][]{cole00,navarro00b,vdbosch00}: it is possible to reproduce the TF relation in one passband easily without reproducing the TF relation in other passbands \\citep[for a multi-waveband comparison of models with the TF relation see e.g.][] {heavens99}.  One way around this confusion is to explore the total baryonic mass TF relation. An estimate of the baryonic TF relation can be obtained by adding the gas mass to a crude estimate of stellar mass implied by the luminosity (usually assuming a constant \\mln). This has been attempted the most thoroughly by \\citet{mcgaugh00} using a constant \\ml in $B$, $I$, $H$, and $K$ bands, although e.g.\\ \\citet{milgrom88} and \\citet{matthews98} discussed aspects of this problem.  Because the stellar \\ml is likely to vary along the TF relation in all passbands, their composite baryonic TF relation will have a larger scatter and different slope than the true baryonic TF relation. A deeper and firmer understanding of the baryonic TF relation is only possible once variations in stellar \\ml along the TF relation are understood and incorporated in the analysis.  In this paper, we use simplified spiral galaxy evolution models similar to the ones presented by \\citet{papiii} to investigate plausible trends in stellar \\ml with galaxy properties, assuming a universal IMF. We discuss these models briefly in \\S \\ref{sec:mod}.  In \\S \\ref{sec:ml} we investigate trends in spiral galaxy stellar \\ml for a number of plausible models, finding that there are systematic variations in stellar \\ml as a function of many galaxy parameters, and that stellar \\mls correlate most tightly with galaxy color. In \\S \\ref{sec:unc} we investigate the physical basis of the color--\\ml relation and we discuss uncertainties in the stellar \\mlsn, including the effects of using different stellar population models, different IMFs, different galaxy evolution prescriptions, and dust. In \\S \\ref{sec:rotn} we discuss the implications of these variations in stellar \\ml for rotation curve decompositions, and put the stellar \\mls onto an observationally-determined maximum-disk scale. In \\S \\ref{sec:tf} we then discuss at length the implications of these variations in stellar \\ml for the stellar mass and baryonic TF relation. Finally, in \\S \\ref{sec:conc}, we present our conclusions. Readers not interested in the details of the models and a detailed analysis of the uncertainties in model stellar \\mls can skip \\S\\S \\ref{sec:mod} and \\ref{sec:unc}. Note that we state all stellar \\mls in solar units. We adopt the HST Key Project distance scale in this paper, corresponding to $H_0 = 71$ km\\,s$^{-1}$\\,Mpc$^{-1}$ \\citep{sakai00}. ",
        "conclusions": "\\label{sec:conc}  Under the assumption of a universal spiral galaxy IMF, we have used stellar population synthesis models in conjunction with simplified spiral galaxy evolution models to argue that there are substantial variations in stellar \\ml in optical and near-IR passbands, and that these \\ml variations are strongly correlated with stellar population colors.  The variations in stellar \\ml also correlate with other galaxy properties (albeit with more scatter) such that, on average, low surface brightness, high gas fraction and low luminosity galaxies have lower stellar \\mls than high surface brightness, low gas fraction, bright galaxies.  The changes in stellar \\mls over a plausible range of galaxy parameters amount to a factor of about 7 in $B$, 3 in $I$, and 2 in $K$ band.  In addition, because the central regions of galaxies are often redder than their outer regions, the inner regions of galaxies are likely to have larger stellar \\mls than the outer regions of galaxies.  This strong correlation between color and stellar \\ml is robust to uncertainties in stellar population and galaxy evolution modeling, including the effects of modest bursts of recent star formation. Larger bursts, which are correspondingly more rare and are typically selected against in spiral galaxy studies (as evidenced by the modest scatter in our TF relation), may depress the stellar \\ml from our expectations by up to 0.5 dex, at most. In addition, because dust both dims and reddens the light from galaxies, uncertainties in the exact amount of dust do not significantly affect the stellar mass estimate for a given galaxy.  The stellar IMF remains the primary uncertainty, implying that these trends are relative in a robust sense, but the absolute normalization is somewhat uncertain.  We analyzed observed $K$ band maximum disk stellar \\mls from \\citet{verheijen97} to place the relative stellar \\mls which we estimate in this paper on the maximum allowed scale.  We find that a Salpeter IMF is ruled out by this analysis if the Ursa Major Cluster is placed at the HST Key Project distance of 20.7 Mpc, and that modification of the low-mass end of the IMF is required. We find that the observed maximum disk stellar \\mls follow the trend suggested by the models, which lends independent support for our models, implies that a fraction of high surface brightness galaxies are reasonably close to maximum disk, and suggests a universal spiral galaxy stellar IMF.  We apply these maximum disk-scaled trends in stellar \\ml with galaxy color to investigate the underlying nature of the TF relation.  We find that, using mass-dependent dust extinction corrections and the color-dependent stellar \\mls it is possible to estimate stellar masses from different passbands which are consistent at better than the 10\\% level. The slope of the stellar mass TF relation of the Ursa Major Cluster sample is 4.4 $\\pm$ 0.2, using an least-squares bisector fit. Including the contribution from the \\hi mass, we find that the slope of the baryonic TF relation of the Ursa Major Cluster is 3.5 $\\pm$ 0.2 (random) $\\pm$ 0.2 (systematic), adopting an unweighted bisector fit.  The slope will be shallower, if all disks are substantially sub-maximal. This is considerably shallower than the baryonic TF of \\citet{mcgaugh00}, who advocate $L \\propto V^{4 \\pm 0.1}$.  We attribute the bulk of the difference to a difference in distance scales and stellar \\ml normalization."
    },
    "0011/astro-ph0011170_arXiv.txt": {
        "abstract": "We present an analysis of the Internal Shock Model of GRBs, where gamma-rays are produced by internal shocks within a relativistic wind. We show that observed GRB characteristics impose stringent constraints on wind and source parameters. We find that a significant fraction, of order 20\\%, of the wind kinetic energy can be converted to radiation, provided the distribution of Lorentz factors within the wind has a large variance and provided the minimum Lorentz factor is $>\\Gamma_{\\pm}\\approx10^{2.5}L_{52}^{2/9}$, where $L=10^{52}L_{52}{\\rm erg\\,s}^{-1}$ is the wind luminosity. For a high, $>10\\%$, efficiency wind, spectral energy breaks in the 0.1 to 1~MeV range are obtained for sources with dynamical time $R/c\\lesssim1$~ms, suggesting a possible explanation for the observed clustering of spectral break energies in this range. The lower limit $\\Gamma_{\\pm}$ to wind Lorenz factor and the upper limit $\\approx1(R/10^7{\\rm cm})^{-5/6}$~MeV to observed break energies are set by Thomson optical depth due to $e^{\\pm}$ pairs produced by synchrotron photons. Natural consequences of the model are absence of bursts with peak emission energy significantly exceeding 1~MeV, and existence of low luminosity bursts with low, 1~keV to 10~keV, break energies. ",
        "introduction": "The widely accepted interpretation of the phenomenology of $\\gamma$-ray bursts (GRBs) is that the observable effects are due to the dissipation of the kinetic energy of a relativistically expanding wind, a ``fireball'' (see \\cite{Meszaros95} and \\cite{Piran96} for reviews). The discovery of GRB afterglow emission in X-ray (Costa et al. 1997), optical (van Paradijs et al. 1997) and radio wavelengths (Frail et al. 1997) confirmed (Waxman 1997, Wijers, Rees \\& M\\'esz\\'aros 1997), standard model predictions of afterglow (Paczy\\'nski \\& Rhoads 1993, Kats 1994, M\\'esz\\'aros \\& Rees 1997, Vietri 1997) that results from the collision of the expanding fireball with surrounding medium.  Both the spectrum and temporal dependence of afterglow emission are consistent with synchrotron emission of electrons accelerated to high energy at the shock wave driven by the fireball into its surrounding medium. The rapid temporal variability of $\\gamma$-ray emission, i.e. of the GRB itself, imply, on the other hand, that the GRB is produced by internal shocks within the expanding wind (\\cite{Woods95,SNP97}; Daigne \\& Mochkovitch, 1998; Panaitescu, Spada \\& M\\'esz\\'aros 1999). Both synchrotron and inverse-Compton emission from shock accelerated electrons have been proposed as the GRB emission mechanism. However, synchrotron emission is favored if the fireball is required to be ``radiatively efficient,'' i.e. if a significant fraction of the fireball energy is required to be converted to $\\gamma$-rays (e.g. \\cite{Derishev00}). Observations of GRB970508 afterglow implies a high radiative efficiency for this GRB (\\cite{FWK00}), and X-ray afterglow observations imply that high radiative efficiency is required for most GRBs (\\cite{Freedman00}).  Observations therefore suggest that the GRB emission is due to synchrotron emission of electrons accelerated by internal fireball shocks. This model faces, however, three major difficulties. First, it is generally argued that if $\\gamma$-ray emission is due to internal shocks within the fireball, then only a small fraction, $<10^{-2}$, of fireball energy is converted to $\\gamma$-rays (Daigne \\& Mochkovitch, 1998; \\cite{Kumar,KnP99,Spada}). This is in conflict with the high radiative efficiency implied by afterglow observations. Second, the $\\gamma$-ray emission of more than 60\\% of the GRBs observed by BATSE detectors is peaked at photon energies in the range of 50~keV to 300~keV (Brainerd et al. 1999). It is generally claimed that clustering of peak emission energy in this narrow energy range requires fine tuning of fireball model parameters, which should naturally produce a much wider range of peak emission energies. Finally, there is evidence that, at least in some cases, the GRB spectra at low energy is steeper than expected for synchrotron emission (\\cite{Preece98,Frontera00}).  The steep low energy spectra has led several authors to consider inverse-Compton emission as the source of $\\gamma$ radiation (e.g. Lazzati, et al 2000). However, steep low energy spectra may also be accounted for within the context of the synchrotron emission hypothesis by the effects of inverse-Compton suppression at low energies (\\cite{Derishev00}) and by the contribution to observed $\\gamma$-ray radiation of photospheric fireball emission (e.g. \\cite{MnR00}). As we show here (see discussion in \\S4), the contribution of latter effect is expected to be significant.  The main goal of the present paper is to determine the constraints imposed by GRB observations on the relativistic fireball wind parameters, and hence on the underlying source producing the wind, under the assumption that the observed radiation is due to synchrotron emission produced in internal wind shocks. In particular, we address the questions of whether, and under what conditions, high radiative efficiency can be obtained, and whether the clustering of peak emission energies can be naturally explained by the model.   It has been demonstrated in earlier work (e.g. \\cite{Kobayashi97}; Beloborodov 2000), that high radiative efficiency may be obtained in  the internal shock model provided the variance of the wind Lorentz factors distribution is large. Kobayashi et al. (1997) found $\\sim10\\%$ efficiency for uniform Lorentz factor distribution with maximum to minimum Lorentz factor ratio $\\approx10$, and Beloborodov (2000) demonstrated that the efficiency may approach 100\\% for non-uniform distributions with wide range of Lorentz factors\\footnote{Kobayashi \\& Sari (2001) similarly find $\\sim60\\%$ efficiency for non-uniform distribution of Lorentz factors over the range $10$ to $10^4$.}. In the analysis presented here we consider several issues which have not been addressed in earlier work: The constraints imposed on the model by the observed peak emission energy distribution, the effect of optical depth due to $e^\\pm$ pairs produced within the fireball wind, and the upper limit imposed on the maximum wind Lorentz factor by the acceleration process. The analysis presented here is also more general, as we study the dependence of radiative efficiency on a wider set of models parameters. We find that the radiative efficiency is limited to values well below the 80\\%--90\\% derived by Beloborodov. The inclusion of $e^\\pm$ pair optical depth effects is essential in determining both the efficiency and the peak emission energy.  High radiative efficiency clearly requires that a large fraction of the internal energy generated by internal shocks be carried by shock accelerated electrons, i.e. that the electron and proton energy densities be close to equipartition. Moreover, high radiative efficiency requires the magnetic field energy density to be not far below equipartition as well (e.g. \\cite{Derishev00}). Thus, we assume in most our calculations that electrons, protons, and magnetic field energy densities are close to equipartition, and focus on the study of the dependence of observed $\\gamma$-ray flux and spectrum on wind luminosity and on wind dynamical parameters (source size, variability time, Lorenz factor and mass distribution of shells within the wind). For completeness, we demonstarte using our calculations (in \\S3.2) that near equipartition is required in order for model predictions to be consistent with the observed peak emission energy distribution.  The model outline is presented in \\S2. The results of a numerical analysis of the model are presented and discussed in \\S3. The main conclusions are summarized in \\S4. ",
        "conclusions": "We have analyzed a model of GRBs, in which a compact source of linear scale $R_0$ produces a wind characterized by an average luminosity $L_w$ and mass loss rate $\\dot M=L_w/\\eta c^2$. At small radius, the wind bulk Lorentz factor, $\\Gamma$, grows linearly with radius, until most of the wind energy is converted to kinetic energy and $\\Gamma$ saturates at $\\Gamma\\sim\\eta$. Variability of the source results in fluctuations in the wind saturation Lorentz factor $\\Gamma$, leading to internal shocks in the expanding wind. These shocks reconvert a fraction of the kinetic energy back to internal energy, which is assumed to be radiated as $\\gamma$-rays by synchrotron (and inverse-Compton) emission of shock-accelerated electrons. Since the wind duration, $t_w\\sim10$~s, is much larger than the dynamical time of the source, $t_d\\sim R_0/c$, variability of the wind on a wide range of time scales, $t_d<t_v<t_w$, is possible. For simplicity, we have assumed that in addition to $t_d$, which determines the initial shell thickness $\\Delta=c t_d\\sim R_0$, and $t_w$, the wind is characterized by a single time scale $t_v>t_d$, which determines the shell ejection rate $t_v^{-1}$. We have addressed the questions of whether, and under what conditions, high radiative efficiency consistent with observations can be obtained, and whether the observed clustering of peak emission energies can be naturally explained by the model.  We have shown that a significant fraction, $\\sim15\\%$, of the fireball energy can be converted to radiation. As pointed out at the end of \\S2, our simplified treatment of post-collision shell evolution, assuming complete merger of shells, leads to an overestimate of the reduction of the variance of colliding shells' Lorenz factors with wind evolution. A more detailed calculation of the merger process will lead to an enhancement of the radiative efficiency by a factor of order unity. The exact value of the enhancement factor will depend, however, on the unknown internal shell structure. Adopting simplifying assumptions regarding the shell structure and post-collision evolution, Kobayashi \\& Sari (2001) have recently pointed out that including post-collsion evolution effects may lead to radiative efficiency exceeding $\\epsilon_e$ (note, that this may be the case also when post-collision evolution is neglected). Although this effect may relax somewhat the constraint imposed on $\\epsilon_e$ by the requirement of high efficiency, we have shown here (see following paragraph), that $\\epsilon_e$ is constrained to be of order unity by the observed $\\gamma$-ray spcetra.  In order to obtain high radiative efficiency and peak emission energy $\\sim1$~MeV, the minimum radius $R_i$ at which internal collisions in the expanding wind occur is required to be similar to $R_\\pm$, the radius where the Thomson optical depth due to $e^\\pm$ pairs produced by shock synchrotron emission equals unity [see Eq. (\\ref{rpmmin})], and a large variance (compared to the mean) of the colliding shells' Lorenz factor (LF) distribution is required. In addition, the electron and magnetic field energy fractions should be close to equipartition, with $\\epsilon_e\\gtrsim0.1$ and $\\epsilon_B\\gtrsim0.01$, in order for model predictions to be consistent with the observed peak emission energy distribution [see Fig. 4 and Eq.(\\ref{eq:emax1})]. The constraint $R_i\\sim R_\\pm$ is equivalent to a constraint on the minimum LF, $\\Gamma_m$, of expanding shells, $\\Gamma_m\\sim\\Gamma_{m\\pm}\\approx10^{2.5}(L_{w,52}/t_{v,-3})^{2/9}$, where $L_w=10^{52}L_{w,52}{\\rm erg\\,s}^{-1}$ and $t_v=10^{-3}t_{v,-3}$~s [see Eq. (\\ref{eq:gpm})]. Large variance in the LF distribution of colliding shells than requires a non-uniform LF distribution (e.g. truncated log-normal or bimodal distributions, see Fig. 1) with $\\Gamma_M$, the maximum LF of wind shells, close to the upper limit set by the shell acceleration process, $\\Gamma<\\eta_*\\approx(\\sigma_T L_w/4\\pi m_p c^3 R_0)^{1/4}= 2\\times10^3(L_{w,52}/R_{0,6})^{1/4}$, where $R_0=10^6R_{0,6}$~cm.  Our results do not agree with the very high, $\\sim80\\%$, radiative efficiency values, obtained in the calculations of Beloborodov (2000) and Kobayashi \\& Sari (2001) by assuming wind Lorentz factor distributions extending from $\\Gamma_m\\sim10$ to $\\Gamma_M\\gtrsim10^4$. In our opinion the minimum wind Lorentz factor must be significantly higher than 10 (as otherwise the wind becomes optically thick) while the maximum Lorentz factor can not significantly exceed $10^3$ (due to the acceleration process limitations).  We have shown that there is an upper limit to the observed energy of photons, $\\varepsilon_p$,  at which the $\\gamma$-ray flux peaks, $\\varepsilon_p\\lesssim\\varepsilon_p^{\\max} \\approx 0.5 R_{0,7}^{-5/6}$~MeV (where $R_0=10^7R_{0,7}$~cm), with very weak dependence on $L_w$ and on $t_v$ [see discussion preceding Eq. (\\ref{eq:emax1}), and figures 2, 5 and 6]. Thus, the source dynamical time $t_d\\sim R_0/c$ must satisfy $t_d\\lesssim 1$~ms in order to allow $\\varepsilon_p\\sim1$~MeV. $\\varepsilon_p\\approx\\varepsilon_p^{\\max}$ is obtained for $\\Gamma_m\\approx\\Gamma_{m\\pm}$, for which the radiative efficiency is largest, and $\\varepsilon_p$ values in the range of 0.1~MeV to 1~MeV are obtained for wind parameters for which the radiative efficiency is high, $\\gtrsim10\\%$ (see figure 3).  High radiative efficiency and peak emission energy consistent with observations are therefore obtained for $\\Gamma_m\\sim\\Gamma_{m\\pm}$. This does not necessarily imply that fine tuning of this model parameter is required. For $\\Gamma_m<\\Gamma_{m\\pm}$, most efficient collisions occur at radii where the optical depth is high, leading to low efficiency, and hence low luminosity, bursts with peak emission energy $\\sim1$~keV (see figures 2 and 3), which would not have been detected by BATSE. For $\\Gamma_m$ significantly higher than $\\Gamma_{m\\pm}$, LF variance is small, leading to low efficiency, low luminosity bursts with peak emission energy $\\sim10$~keV, which may have been difficult to detect with BATSE. Natural consequences of the model considered here are therefore absence of bursts with peak emission energy significantly exceeding $\\sim1$~MeV, and existence of low luminosity bursts with low, $\\sim1$~keV to $\\sim10$~keV, peak emission energy. The frequency of such bursts depends on the distribution of $\\Gamma_m$ in different winds.  It should be pointed out in this context that while the deficit, among bursts detected by BATSE, of bursts with peak emission energy below $\\sim50$~keV clearly reflects a real deficit of such bursts, the deficit in detected bursts with peak emission energy $>0.5$~MeV may be partly due to a selection effect (e.g. \\cite{Lloyd99}). BATSE data are consistent with equal number of bursts per logarithimic peak energy interval beyond $\\varepsilon_p\\sim0.5$~MeV (with luminosity comparable to that of lower peak energy bursts). Our model prediction, that bursts with $\\varepsilon_p\\gg1$~MeV should be absent, should therefore be tested with future GRB detectors, which are able to better constrain the high energy end of the $\\varepsilon_p$ distribution.  A note should be made here regarding the low energy spectral slope. The upper limit $\\eta_*$ on $\\Gamma_M$ is a consequence of the fact that for $\\eta>\\eta_*$ shell acceleration saturates at $\\Gamma\\sim\\eta_*<\\eta$, as the wind becomes optically thin to Thomson scattering, and the internal energy left in the shell escapes as thermal radiation rather than being converted to kinetic energy. The requirement $\\Gamma\\sim\\eta_*$ therefore implies that a significant fraction of the wind energy may escape as thermal radiation, leading to low energy spectral slopes steeper than those expected for pure synchrotron emission. This may account, at list partially, for observed steep low energy spectra.  Finally we note that the lower limit imposed on $\\Gamma_m$, $\\Gamma_m\\gtrsim\\Gamma_\\pm$, is not derived from the requirement that the pair production optical depth for high energy, $>100$~MeV, photons be smaller than unity. While this requirement leads to a similar constraint, $\\Gamma_m\\gtrsim10^2$, high energy photons have been detected in a small number of cases only. The constraint $\\Gamma_m>\\Gamma_{m\\pm}$ is imposed in the present analysis by the requirement that the wind Thomson optical depth due to $e^\\pm$ pairs produced by synchrotron photons be smaller than unity at the internal shocks stage."
    },
    "0011/astro-ph0011200_arXiv.txt": {
        "abstract": "We demonstrate the capability of Principal Component Analysis (PCA) as applied by Heyer \\& Schloerb (1997) to extract the statistics of turbulent interstellar velocity fields as measured by the energy spectrum, $E(k) \\sim k^{-\\beta}$.  Turbulent velocity and density fields with known statistics are generated from fBm simulations.  These fields are translated to the observational domain, T(x,y,v), considering the excitation of molecular rotational energy levels and radiative transfer.  Using PCA and the characterization of velocity and spatial scales from the eigenvectors and eigenimages respectively, a relationship is identified which describes the magnitude of line profile differences, ${\\delta}v$ and the scale, $L$, over which these differences occur, ${\\delta}v \\sim L^\\alpha$. From a series of models with varying values of $\\beta$,  we find, $\\alpha = 0.33\\beta - 0.05$ for $1 < \\beta < 3$. This provides the basic calibration between the intrinsic velocity field statistics to observational measures and a diagnostic for turbulent flows in the interstellar medium.  We also investigate the effects of noise, line opacity, and finite resolution on these results. ",
        "introduction": "The evolution of the cold, dense interstellar medium and the formation of stars are regulated by the kinetic energy content of turbulent gas flows. Such chaotic motions play a critical role in the equilibrium of the gas as these provide a countering pressure to self gravity and the pressure of the external medium. In the absence of external pressure fronts, the formation of stars occurs exclusively in dense, localized regions where the turbulent energy is sufficiently reduced to enable gravitational collapse. Therefore, an understanding of the dense interstellar medium and the formation of stars  requires an accurate description of turbulent gas flow including energy injection into and dissipation from the gas system and the distribution of turbulent kinetic energy over varying spatial scales.  Turbulent flows within the molecular interstellar medium are recognized as observed supersonic line widths and morphological complexity (Scalo 1990; Falgarone Phillips, \\& Walker 1991).  While these observations denote the presence of turbulence in the interstellar medium, there is little quantitative information as to the nature of the turbulent flows. Ideally, one would like to recover the 3-dimensional velocity field, $v_z(x,y,z)$, for the line of sight component, z, from a set of spectroscopic observations.  In practice, the reconstruction of the true velocity field from observations is not possible given the chaotic, non-deterministic nature of turbulent flows and the non-linear transformation of the field onto the spectroscopic axis of the observer. Therefore, it is essential to construct statistical descriptions of turbulence (Dickman 1985). The primary statistic used to distinguish turbulent models is the energy spectrum, E(k), which describes the degree of coherence of the velocity field over spatial scales. In D dimensions,  the energy spectrum is the angular integral of the power spectrum, $$ {E(k) = \\int d\\Omega P(k)  \\propto <P(|k|)>_{\\Omega} k^{D-1}} \\eqno(1.1) $$ We refer to the spectral index $\\beta$ in connection with E(k), $$ {E(k) \\sim |k|^{-\\beta} }  \\eqno(1.2) $$ The power spectrum then has the form (in D dimensions) : $$ {P(k) \\sim |k|^{-(\\beta+(D-1))} } \\eqno(1.3) $$ For a dissipationless cascade of energy through an incompressible fluid, $\\beta = 5/3$ and the power spectrum, P(k), decreases in 3 dimensions as k$^{-11/3}$ (Kolmogorov 1941). A 2 or 1 dimension cut through the 3 dimensional structure will also have an energy spectrum with the same index $\\beta$ but with a different slope of the power spectrum.  The mean square turbulent velocity, $v_l$,  at size $l$, is determined from the spectrum of velocity fluctuations over scales smaller than $l$, $$ v_l^2 \\sim \\int_{2\\pi/l}^\\infty E(k)dk \\propto \\int_{2\\pi/l}^\\infty k^{-\\beta}dk \\propto l^{\\beta-1} \\eqno(1.4)$$ $v_l^2$ describes the variance of the 3 dimensional velocity field, v(x,y,z) within a volume $l^3$. The root mean square turbulent velocity, $<v_l^2>^{1/2}$, is then, $$  <v_l^2>^{1/2} \\propto l^{(\\beta-1)/2} = l^\\gamma \\eqno(1.5) $$ where $\\gamma=(\\beta-1)/2$. It is not equivalent  to an observed width of a spectral line which depends upon the velocity, density, temperature, and abundance which vary along the line of sight. However, given a velocity field described by an energy spectrum with $\\beta > 0$, one might expect  some correlation of measured velocities, ${\\delta}v$, with projected spatial scale, L, of the form, $${\\delta}v \\sim L^\\alpha \\eqno(1.6)$$ where $\\alpha$ refers to the measured spectral index in contrast to the intrinsic index $\\gamma$ in equation (1.5). Conventionally, ${\\delta}v$ has been derived from measured line width or the displacement of centroid velocities of spectral lines.  Coarse limits to values of $\\beta$ appropriate for the molecular interstellar medium  can be inferred from the early models of observed CO line profiles. To account for the broad velocity widths and the ability of the high opacity CO lines to detect the warm molecular gas located deep within the Orion cloud, Goldreich \\& Kwan (1974) modeled line profiles with a purely systematic velocity field replicating collapsing motions.  For such smooth, systematic velocity fields, $\\beta \\approx 3$.   However, it is unlikely that all molecular clouds undergo global gravitational collapse since this predicts too large a star formation rate in the Galaxy (Zuckerman \\& Evans 1974). Other systematic velocity fields such as rotation are not observed over the extended scales of molecular clouds (Arquilla \\& Goldsmith 1984). These results suggest $\\beta$ must be smaller than 3. A lower limit to $\\beta$ is provided by the inability of purely random velocity fields ($\\beta $ = -2 in 3D) to reproduce observed line profiles (Leung \\& Lizst 1976; Dickman 1985). Given these considerations, an approximate range to the spectral index of the energy spectrum is $ 3/2 < \\beta < 5/2$.  More direct measures of $\\beta$ rely on the identification of a size-line width relationship and the assumption that variations along the spectroscopic axis accurately reflect variations within the intrinsic velocity field. Using a set of inhomogeneous data from the literature, Larson (1981) determined a relationship between the global velocity dispersion, $\\sigma_v$, and size, L, for a sample of 46 molecular regions in the solar neighborhood, $\\sigma_v \\sim L^{0.38}$.  Using a more homogeneous sample of observations, Solomon etal (1987) identified a relationship with an index of 0.5.  However, the applicability of these observed relationships to the energy spectrum associated with interstellar turbulence (equation (1.5))  is weak unless one considers the targeted clouds as part of singular, global fluid within the Galaxy (see Dickman 1985). Given the wide range of distances and location of the clouds within different spiral arms, such an assumption is clearly invalid. Rather, these respective relationships derived from a sample of unrelated clouds provide a statement of self gravitational equilibrium if one assumes that the mean column density of clouds is approximately constant.   That is, there is a sufficient amount of mass distributed within the cloud to bind the internal turbulent motions.  Several investigations have decomposed spectroscopic data cubes of \\coa~ or \\cob~ emission from targeted molecular clouds into discrete units or ``clumps'' based on some preassigned definition (Carr 1987; Stutzki \\& Gusten 1990; Williams, de Geus, \\& Blitz 1994). For each clump, a global velocity dispersion and size are determined. However, the resultant size-line width relationships show a large scatter and little significant correlation.  This may reflect that the localization of emission within a spectroscopic data cube may not uniquely originate from a discrete clump localized in  space.  Also, the process of identifying clumps limits the resolution to the size of the smallest object.  Scalo (1984) and Kleiner \\& Dickman (1985) proposed the use of structure and autocorrelation functions of centroid velocity images derived from spectroscopic data cubes.  This assumes that the 2 dimensional projection retains the statistical characteristics of the 3 dimensional velocity field. These correlation methods have been applied by Miesch \\& Bally (1994) to a sample of 12 molecular regions imaged in \\coa~ emission.  The structure functions are described by  a power law with a spectral index of 0.86 which corresponds to $\\alpha$=0.43 as in equation (1.6). However, these correlation methods are limited. For many regions in the molecular ISM, observed line profiles exhibit non-gaussian shapes for which the centroid velocity is not an accurate approximation.  Also, the centroid velocity is not well defined near the edges of the clouds where the signal is weak. Finally, the relationship between the measured value of $\\alpha$ and the intrinsic velocity field statistics, $\\gamma$ or $ \\beta$, has not yet been established.  Recently, Heyer \\& Schloerb (1997, hereafter HS97)  motivated the use the multivariate technique of Principal Component Analysis (PCA) to decompose spectral line imaging observations of molecular clouds. From the analysis, they identify a relationship which describes the magnitude of velocity differences of line profiles and the spatial scale at which these differences occur within the image. For several molecular clouds, they find $\\alpha \\sim 0.4 - 0.5$. However, it has not yet been demonstrated that this analysis, or the other methods described above, are sensitive to varying field statistics ($\\beta$) nor has the relationship between $\\alpha$ and $\\beta$ been established. Given the complex transformation of the velocity field to the observed spectral axis, one can not assume that $\\alpha = \\gamma = (\\beta-1)/2 $.  In this paper, we quantitatively evaluate the ability of PCA to characterize intrinsic velocity fields from spectral line observations. Guided by theoretical and observational results, 3 dimensional  models of velocity and density fields with known statistical properties are generated. Observations of these fields are simulated, accounting for non-LTE excitation and radiative transfer.  The simulated observations are analyzed with PCA under a variety of physical conditions and observational limitations. We demonstrate that the method is indeed sensitive to velocity fields with varying energy spectra and derive the appropriate calibration between the observed quantity $\\alpha$ and $\\beta$. In a companion paper (Brunt \\& Heyer 2000), we  use these results to derive the energy spectra for molecular regions in the outer Galaxy. ",
        "conclusions": "We have investigated the ability of PCA to recover intrinsic turbulent velocity field statistics from spectral line imaging observations of the molecular ISM, and established a preliminary calibration of the method. The calibration, subject to the limitations of the modeling, identifies a monotonic relationship between the intrinsic $\\beta$ values (or equivalently, $\\gamma$) and observed exponents ($\\alpha$) that characterize the velocity field statistics.  In addition to the simplistic nature of the input density and velocity fields, including the assumption of statistical isotropy and the periodicity feature, the major limitations are the lack of mutual consistency between density and velocity, and the assumption of local excitation in the radiative transfer calculations. However, the simulated lognormal observations are in good visual correspondence with real molecular ISM observations. The retrievals from the low density observations are in close agreement with those from the high density observations. This suggests that the accuracy of the method depends upon the degree to which one samples the velocity field, rather than the details of the radiative transfer. For molecular clouds in particular, highly saturated observations may be more reliable (due to more homogeneous sampling and a greater number of retrieved measurements) and this suggests that $^{12}$CO is likely to provide the most accurate values of $\\alpha$. This method could also be applied to HI 21cm emission which is excited over a wide range of physical conditions.  The sampling effects due to noise, which uniformly masks sources of variance, do not affect the fits for $\\alpha$, until very high noise levels are reached. As calibrated, the PCA method is resolution-independent but limited in general by necessary assumptions of sub-pixel homogeneity of the emission field. Independent of our modeled calibration, the robustness of measured $\\alpha$ values provide quantitative constraints to simulated observations of velocity and density fields from hydrodynamical calculations. We do stress, however, that the calibration of $\\alpha$ to $\\beta$ achieved in this paper is strictly only valid for the fBm model velocity fields. More realistic turbulent fields may not adhere as closely to the calibration, and in future we intend to include supercomputer MHD simulations in the PCA calibration.  Empirical measurements of $\\alpha$ from real data of course do not depend on the calibration to $\\beta$ (or $\\gamma$), and such measurements will be reported by Brunt \\& Heyer (2000) for an ensemble of fields taken from the FCRAO Outer Galaxy Survey. Using the preliminary calibration established here, we obtain constraints on the intrinsic velocity field statistics of the molecular ISM.  This work was supported by NSF grant AST 97-25951 to the Five College Radio Astronomy Observatory.  \\clearpage"
    },
    "0011/astro-ph0011036_arXiv.txt": {
        "abstract": "s{ We analyze the effect of primordial density perturbations      % on the cosmic QCD phase transition. According to our results   % hadron bubbles nucleate at the cold perturbations.             % We call this mechanism inhomogeneous nucleation.               % We find the typical distance between bubble centers to be a few meters. This exceeds the estimates from homogeneous nucleation by two orders of magnitude. The resulting baryon inhomogeneities may affect primordial nucleosynthesis. }  \\clearpage  \\pagestyle{plain}    The order of the QCD transition and the values of its parameters are still under debate. Nevertheless there are indications from lattice QCD calculations. Quenched QCD (no dynamical quarks) shows a first-order phase transition with a small latent heat, compared to the bag model, and a small surface tension, compared to dimensional arguments \\cite{Iwasaki}. We assume that the QCD transition is of first order and that the values from quenched lattice QCD (scaled appropriately by the number of degrees of freedom) are typical for the physical QCD transition. Based on these values and homogeneous bubble nucleation a small supercooling, $\\Delta_{\\rm sc} \\equiv 1- T_{\\rm f}/T_{\\rm c} \\sim 10^{-4}$, and a tiny bubble nucleation distance, $d_{\\rm nuc} \\sim 1$ cm, would    % follow \\cite{Ignatius}. The actual nucleation temperature is denoted by $T_{\\rm f}$, and the thermodynamic transition temperature by                 % $T_c \\approx 150$~MeV.                                          %  We argue~\\cite{IS}                                              % that the assumption of homogeneous nucleation is violated in the early Universe by the inevitable density perturbations from inflation or from other seeds for structure formation. Those fluctuations in density and temperature have been measured by COBE \\cite{Bennett} to have an amplitude of $\\delta T/T \\sim 10^{-5}$. The effect of the QCD transition on density perturbations \\cite{SSW,SSW2} and gravitational waves \\cite{Schwarz} has been studied previously, while we investigate the effect of the density perturbations on the QCD phase transition here.  First-order phase transitions normally proceed via nucleation of bubbles of the new phase. When the temperature is spatially uniform and no significant impurities are present, the mechanism is homogeneous nucleation. The probability to nucleate a bubble of the new phase per time and volume is approximated by $\\Gamma \\approx T_c^4 \\exp[-S(T)]$. The nucleation action $S$ is the free energy difference of the system with and without the nucleating bubble, divided by the temperature.  Nucleation is a very rapid process, compared with the extremely slow cooling of the Universe. The duration of the nucleation period, $\\Delta t_{\\rm nuc}$, is found to be \\cite{Fuller,Enqvist} \\begin{equation} \\Delta t_{\\rm nuc} = - \\frac{ \\pi^{1/3} }{{\\rm d}S / {\\rm d}t \\left|_{t_f} \\right.} . \\end{equation} The time $t_f$ is defined as the moment when the fraction of space where nucleations still continue equals $1/e$. The heat flow preceding the deflagration fronts reheats the rest of the Universe. We denote by $v_{\\rm heat}$ the effective speed by which released latent heat propagates in sufficient amounts to shut down nucleations. In practice, $v_{\\rm def} < v_{\\rm heat} < c_s$, where $v_{\\rm def}$ is the velocity of the deflagration front and $c_s$ is the sound speed \\cite{Kurki-Suonio}. In the unlikely case of detonations $v_{\\rm heat}$ should be replaced by the velocity of the phase boundary in all expressions that follow.  The mean distance between nucleation centers, measured immediately after the transition completed, is \\begin{equation} d_{\\rm nuc,hom} = 2 v_{\\rm heat} \\Delta t_{\\rm nuc} \\label{dnuceq}. \\end{equation} This nucleation distance sets the spatial scale for baryon number inhomogeneities.  Lattice simulations imply that in real-world QCD the energy density must change very rapidly in a narrow temperature interval. This can be seen from the microscopic sound speed in the quark phase, $c_s \\equiv (\\partial p/\\partial \\varepsilon)_{\\rm S}^{1/2}$. Lattice QCD indicates that $3 c_{\\rm s}^2(T_{\\rm c}) = {\\cal O}(0.1)$ \\cite{Latticecs2}. Thus, the cosmological time-temperature relation is strongly modified already before the nucleations, due to \\begin{equation} \\frac{ {\\rm d}T }{ {\\rm d}t } = - 3 c_{\\rm s}^2 \\frac{T}{t_{\\rm H}}, \\label{Tteq} \\end{equation} where $t_H  \\equiv  1/H = (3 M_{\\rm pl}^2/8\\pi\\varepsilon_{\\rm q})^{1/2}$ with $\\varepsilon_{\\rm q}$ being the energy density in the quark phase. This behavior of the sound speed increases the nucleation distance because of the proportionality $\\Delta t_{\\rm nuc}\\propto 1/[3 c_{\\rm s}^2 (T_f)]$ \\cite{Ignatius}.  In the thin-wall approximation the nucleation action has the following explicit expression: \\begin{equation} \\label{C} S(T)  =  \\frac{ C^2 }{(1-T/T_c)^2}, \\,\\,\\,\\,\\,\\, C  \\equiv  4\\sqrt{\\frac{\\pi}{3}} \\frac{\\sigma^{3/2}}{l\\sqrt{T_c}}\\ , \\end{equation} for small supercooling. Assuming further that $c_{\\rm s}$ does not change very much during supercooling, the following relation holds for the supercooling and nucleation scales: \\begin{equation} \\frac{\\Delta t_{\\rm sc}}{\\Delta t_{\\rm nuc}} = \\frac{\\Delta_{\\rm sc}}{\\Delta_{\\rm nuc}} = \\frac{2}{\\pi^{1/3}} \\bar{S} \\label{Seq}. \\end{equation} Here we denote by $\\Delta$ a relative (dimensionless) temperature interval and by $\\Delta t$ a dimensionful time interval. $\\bar{S} \\equiv S(T_f)$ is the critical nucleation action, $\\bar{S} = {\\cal O}(100)$.  Surface tension and latent heat are provided by lattice simulations with quenched QCD only, giving the values $\\sigma = 0.015 T_c^3$, $l = 1.4 T_c^4$~\\cite{Iwasaki}. Scaling the latent heat for the physical QCD leads us to take $l = 3 T_c^4$.  With these values for the latent heat and surface tension, the amount of supercooling is $\\Delta_{\\rm sc} = 2.3 \\times 10^{-4}$. {}From Eq.~(\\ref{Seq}) it follows that $\\Delta_{\\rm nuc} = 1.5 \\times 10^{-6}$. Substituting $3 c_{\\rm s}^2 = 0.1$ into Eq.~(\\ref{Tteq}), we find $\\Delta t_{\\rm nuc} = 1.5 \\times 10^{-5} t_{\\rm H}$ for the duration of the nucleation period. The nucleation distance depends on the unknown velocity $v_{\\rm heat}$ in Eq.~(\\ref{dnuceq}). With the value 0.1 for $v_{\\rm heat}$, the nucleation distance $d_{\\rm nuc,hom}$ would have the value $2.9 \\times 10^{-6} d_{\\rm H}$. One should take these values with caution, due to large uncertainties in $l$ and $\\sigma$. As our reference set of parameters, we take: $\\Delta_{\\rm sc} = 10^{-4}$, $\\Delta_{\\rm nuc} = 10^{-6}$, $\\Delta t_{\\rm nuc} = 10^{-5} t_H$.  In the real Universe the local temperature of the radiation fluid fluctuates. We decompose the local temperature $T(t,{\\bf x})$ into the mean temperature $\\bar{T}(t)$ and the perturbation $\\delta T(t,{\\bf x})$. The temperature contrast is denoted by $\\Delta \\equiv \\delta T/ \\bar{T}$. On subhorizon scales in the radiation dominated epoch, each Fourier coefficient $\\Delta(t,k)$ oscillates with constant amplitude, which we denote by $\\Delta_T(k)$. Inflation predicts a Gaussian distribution, \\begin{equation} \\label{pd} p(\\Delta){\\rm d}\\Delta = {1\\over \\sqrt{2\\pi}\\Delta_T^{\\rm rms}} \\exp\\left( - \\frac12 {\\Delta^2\\over (\\Delta_T^{\\rm rms})^2}\\right) {\\rm d}\\Delta \\ . \\end{equation} We find \\cite{norm} for the COBE normalized \\cite{Bennett} rms temperature fluctuation of the radiation fluid (not of cold dark matter) $\\Delta_T^{\\rm rms} = 1.0 \\times 10^{-4}$ for a primordial Harrison-Zel'dovich spectrum. The change of the equation of state prior to the QCD transition modifies the temperature-energy density relation, $\\Delta = c_s^2 \\delta\\varepsilon /(\\varepsilon + p)$. We may neglect the pressure $p$ near the critical temperature since $p \\ll \\varepsilon_{\\rm q}$ at $T_{\\rm c}$. On the other hand the drop of the sound speed enhances the amplitude of the density fluctuations proportional to $c_s^{-1/2}$ \\cite{SSW2}. Putting all those effects together and allowing for a tilt in the power spectrum, the COBE normalized rms temperature fluctuation reads \\begin{equation} \\Delta_T^{\\rm rms} \\approx 10^{-4} (3 c_s^2)^{3/4} \\left({k\\over k_0}\\right)^{(n-1)/2} \\ , \\end{equation} where $k_0 = (aH)_0$. For a Harrison-Zel'dovich spectrum ($n=1$) and $3 c_s^2 = 0.1$, we find $\\Delta_T^{\\rm rms} \\approx 2 \\times 10^{-5}$.  A small scale cut-off in the spectrum of primordial temperature fluctuations comes from collisional damping by neutrinos \\cite{Weinberg,SSW2}. The interaction rate of neutrinos is $\\sim G_{\\rm F}^2 T^5$. This has to be compared with the angular frequency $c_s k_{\\rm ph}$ of the acoustic oscillations. At the QCD transition neutrinos travel freely on scales $l_{\\nu} \\approx 4 \\times 10^{-6} d_{\\rm H}$. Fluctuations below the diffusion scale of neutrinos are washed out, \\begin{equation} \\label{diff} l_{\\rm diff} = \\left[\\int^{t_c}_0\\!\\! l_{\\nu}(\\bar t) {\\rm d} \\bar t \\right]^{\\frac{1}{2}}\\!\\! \\approx 7\\!\\times\\! 10^{-4} d_{\\rm H} \\label{nueq}\\, . \\end{equation} In Ref.~\\cite{SSW2} the damping scale from collisional damping by neutrinos has been calculated to be $k^{\\rm ph}_{\\nu} = 10^4 H$ at $T=150$ MeV. The estimate (\\ref{diff}) is consistent with this damping scale. We assume $l_{\\rm smooth} = 10^{-4} d_{\\rm H}$. The compression timescale for a homogeneous volume $\\sim l_{\\rm smooth}^3$ is $\\delta t = \\pi l_{\\rm smooth}/c_s \\sim 10^{-3} t_{\\rm H}$. Since $\\delta t \\gg \\Delta t_{\\rm nuc}$ the temperature fluctuations are frozen with respect to the time scale of nucleations. As long as $l_{\\rm smooth}$ exceeds the Fermi scale homogeneous bubble nucleation applies within these small homogeneous volumes. This is a crucial difference to the scenario of heterogeneous nucleation \\cite{Christiansen}, where bubbles nucleate at ad hoc impurities.  Let us now investigate bubble nucleation in a Universe with spatially inhomogeneous temperature distribution. Bubble nucleation effectively takes place while the temperature drops by the tiny amount $\\Delta_{\\rm nuc}$. To determine the mechanism of nucleation, we compare $\\Delta_{\\rm nuc}$ with the rms temperature fluctuation $\\Delta_T^{\\rm rms}$:  1. If $\\Delta_T^{\\rm rms} < \\Delta_{\\rm nuc}$, the probability to nucleate a bubble at a given time is {\\em homogeneous} in space. This is the case of homogeneous nucleation.  2. If $\\Delta_T^{\\rm rms} > \\Delta_{\\rm nuc}$, the probability to nucleate a bubble at a given time is {\\em inhomogeneous} in space. We call this inhomogeneous nucleation.  \\begin{figure}[t] \\begin{center} \\epsfig{figure=itfig1.eps,width=0.325\\linewidth}          % \\epsfig{figure=itfig2.eps,width=0.325\\linewidth}          % \\epsfig{figure=itfig3.eps,width=0.325\\linewidth}          % \\end{center} \\caption{\\label{fig1} Sketch of a first-order QCD transition in the inhomogeneous Universe. At $t_1$ the first hadronic bubbles (H) nucleate at the coldest spots (light gray), while most of the Universe remains in the quark phase (Q). At $t_2$ the bubbles inside the cold spots have merged and have grown to bubbles as large as the temperature fluctuation scale. At $t_3$ the transition is almost finished. The last quark droplets are found in the hottest spots (dark gray).} \\end{figure} The quenched lattice QCD data and a COBE normalized flat spectrum lead to the values $\\Delta_{\\rm nuc} \\sim 10^{-6}$ and $\\Delta_T^{\\rm rms} \\sim 10^{-5}$. We conclude that the cosmological QCD transition may proceed via inhomogeneous nucleation. A sketch of inhomogeneous nucleation is shown in Fig.~\\ref{fig1}. The basic idea is that temperature inhomogeneities determine the location of bubble nucleation. Bubbles nucleate first in the cold regions.  The temperature change at a given point is governed by the Hubble expansion and by the temperature fluctuations. For the fastest changing fluctuations, with angular frequency $c_s/l_{\\rm smooth}$, we find \\begin{equation} \\label{coolingrate} {{\\rm d}T(t,{\\bf x})\\over {\\rm d} t} = {\\bar T\\over t_{\\rm H}} \\left[- 3 c_s^2 + {\\cal O}\\left(\\Delta_T {t_{\\rm H}\\over \\delta t}\\right)\\right]\\ . \\end{equation} The Hubble expansion is the dominant contribution, as typical values are $3 c_s^2 = 0.1$ from quenched lattice QCD and $\\Delta_T^{\\rm rms} t_{\\rm H}/\\delta t \\approx 0.01$ from the discussion above. This means that the local temperature does never increase, except by the released latent heat during bubble growth.  To gain some insight in the physics of inhomogeneous nucleation, let us first inspect a simplified case. We have some randomly distributed cold spheres of diameter $l_{\\rm smooth}$ with equal and uniform temperature, which is by the amount $\\Delta_T^{\\rm rms}T_c$ smaller than the again uniform temperature in the rest of the Universe. When the temperature in the cold spots has dropped to $T_{\\rm f}$, homogeneous nucleation takes place in them. Due to the Hubble expansion the rest of the Universe would need the time $\\Delta t_{\\rm cool} = t_H \\Delta^{\\rm rms}_T  / 3 c_s^2$          % to cool down to $T_{\\rm f}$. Inside each cold spot there is a large number of tiny hadron bubbles, assumed to grow as deflagrations. They merge within $\\Delta t_{\\rm cool}$ if $\\Delta_{\\rm nuc} < (v_{\\rm def}/v_{\\rm heat}) \\Delta^{\\rm rms}_T$. This condition should be clearly fulfilled for our reference set of parameters. Thus the cold spots have fully been transformed into the hadron phase while the rest of the Universe still is in the quark phase. The latent heat released in a cold spot propagates in all directions, which provides the length scale \\begin{equation} \\label{lheat} l_{\\rm heat} \\equiv 2 v_{\\rm heat} \\Delta t_{\\rm cool}. \\end{equation} If the typical distance from the boundary of a cold spot to the boundary of a neighboring cold spot is less than $l_{\\rm heat}$, then no hadronic bubbles can nucleate in the intervening space. In this case the nucleation process is totally dominated by the cold spots, and the average distance between their centers gives the spatial scale for the resulting inhomogeneities. In the following analysis for a more realistic scenario we concentrate in this case, $l_{\\rm heat} > l_{\\rm smooth}$.  The real Universe consists of smooth patches of typical linear size $l_{\\rm smooth}$, their temperatures given by the distribution~(\\ref{pd}). As discussed above, the merging of tiny bubbles within a cold spot can here be treated as an instantaneous process. The fraction of space that is not reheated by the released latent heat (and not transformed to hadron phase), is given at time $t$ by \\begin{equation} \\label{fihn} f(t) \\approx 1 - \\int_0^{t}\\Gamma_{\\rm ihn}(t') V(t,t') {\\rm d}t', \\end{equation} where we neglect overlap and merging of heat fronts. At time $t$ heat, coming from a cold spot which was transformed into hadron phase at time $t'$, occupies the volume $V(t,t') = (4\\pi/3) [l_{\\rm smooth}/2 + v_{\\rm heat}(t-t')]^3$. The other factor in Eq.~(\\ref{fihn}), $\\Gamma_{\\rm ihn}$, is the volume fraction converted into the new phase, per physical time and volume as a function of the mean temperature $T=\\bar{T}(t)$. $\\Gamma_{\\rm ihn}$ is proportional to the fraction of space for which temperature is in the interval $[T_{\\rm f}, T_f(1 + d\\Delta)]$. This fraction of space is given by Eq.~(\\ref{pd}) with $\\Delta = T_{\\rm f}/T - 1$. Rewriting ${\\rm d}\\Delta$ by means of Eq.~(\\ref{Tteq}) leads to the expression \\begin{equation} \\label{gammai} \\Gamma_{\\rm ihn} = 3 c_s^2{T_f\\over T}{1\\over t_{\\rm H} {\\cal V}_{\\rm smooth}} p(\\Delta = \\frac{T_{\\rm f}}{T} - 1) , \\end{equation} where the relevant physical volume is ${\\cal V}_{\\rm smooth} = (4\\pi/3) (l_{\\rm smooth}/2)^3.$  The end of the nucleation period, $t_{\\rm ihn}$, is defined through the condition $f(t_{\\rm ihn}) = 0$. We introduce the variables $N \\equiv (1 - T_{\\rm f}/T)/\\Delta_T^{\\rm rms}$ and ${\\cal N} \\equiv N(t_{\\rm ihn})$. Since $c_s$ may be assumed to be constant during the tiny temperature interval where nucleations actually take place, we find from Eq.~(\\ref{Tteq}): $1 - t/t_{\\rm ihn} \\approx 2/(3c_s^2) \\Delta_T^{\\rm rms} (N - {\\cal N})$. Putting everything together we determine ${\\cal N}$ from \\begin{equation} {l_{\\rm heat}^3\\over l_{\\rm smooth}^3} \\int_{\\cal N}^\\infty {\\rm d} N \\frac{e^{-\\frac12 N^2}}{\\sqrt{2\\pi}} \\left(\\frac{l_{\\rm smooth}}{l_{\\rm heat}} + N - {\\cal N}\\right)^3 = 1 . \\end{equation} The expression                                           % $l_{\\rm heat}/l_{\\rm smooth} =                           % 2 v_{\\rm heat}(3 c_s^2)^{-1/4} (k/k_0)^{(n-1)/2}$        % is valid for the COBE normalized spectrum.               % For $l_{\\rm heat}/l_{\\rm smooth} = 1, 2, 5, 10$ we find ${\\cal N} \\approx 0.8, 1.4, 2.1, 2.6$, respectively.  The effective nucleation distance in inhomogeneous nucleation is defined from the number density of those cold spots that acted as nucleation centers, $d_{\\rm nuc,ihn} \\equiv n^{- 1/3}$. We find \\begin{eqnarray} d_{\\rm nuc,ihn} & \\approx & \\left[\\int_0^{t_{\\rm ihn}} \\Gamma_{\\rm ihn}(t) {\\rm d} t\\right]^{-1/3} \\\\ & = & [\\frac{3}{\\pi} (1 - {\\rm erf}({\\cal N}/\\sqrt{2})]^{-1/3} \\, l_{\\rm smooth}. \\end{eqnarray} With the above values $l_{\\rm heat}/l_{\\rm smooth} = 1, 2, 5, 10$ we get $ d_{\\rm nuc,ihn} = 1.4, 1.8, 3.0, $ $            % 4.8 \\times l_{\\rm smooth}$, where $l_{\\rm smooth} \\approx 1 \\mbox{\\ m}$.  For a COBE normalized spectrum without any tilt and with a tilt of $n-1 = 0.2$ (where $(k_{\\rm smooth}/k_0)^{0.1} \\approx 25$), together with $3c_s^2 = 0.1$ and $v_{\\rm heat} = 0.1$, we find the estimate $l_{\\rm heat}/l_{\\rm smooth} \\approx 0.4$ and 9, correspondingly. Notice that the values of $v_{\\rm heat}$ and $3 c_s^2$ are in principle unknown. Anyway, we can conclude that the case $l_{\\rm heat} > l_{\\rm smooth}$           % is a realistic possibility.  With $2 v_{\\rm heat}(3c_s^2)^{-1/4} (10^{-4}d_{\\rm H}/l_{\\rm smooth}) < 1$ and without positive tilt we are in the region $l_{\\rm heat} < l_{\\rm smooth}$, where the geometry is more complicated and the above quantitative analysis does not apply. In this situation nucleations take place in the most common cold spots (${\\cal N} \\sim 1$), which are very close to each other. We expect a structure of interconnected baryon-depleted and baryon-enriched layers with typical surface $l_{\\rm smooth}^2$ and thickness $l_{\\rm def} \\equiv v_{\\rm def}\\Delta t_{\\rm cool}$. In between $d_{\\rm nuc,hom}$ would be the relevant length scale of inhomogeneities. An accurate analysis of this case requires computer simulations, which is beyond the scope of the present work. However, it is clear that the result will be different compared with homogeneous nucleation.  We emphasize that inhomogeneous and heterogeneous nucleation~\\cite{Christiansen} are genuinely different mechanisms, although they give the same typical scale of a few meters by chance. If latent heat and surface tension of QCD would turn out to reduce $\\Delta_{\\rm sc}$ to, e.g., $10^{-6}$, instead of $10^{-4}$, the maximal heterogeneous nucleation distance would fall to the centimeter scale, whereas on the distance in inhomogeneous nucleation this would have no effect.  We have shown that inhomogeneous nucleation during the QCD transition can give rise to an inhomogeneity scale exceeding the proton diffusion scale (2~m at 150~MeV).                            % The resulting baryon inhomogeneities could provide inhomogeneous initial conditions for nucleosynthesis. Observable deviations from the element abundances predicted by homogeneous nucleosynthesis seem to be possible in that case~\\cite{MathewsKainulainen}.              %  In conclusion, we found that inhomogeneous nucleation leads to nucleation distances that exceed by two orders of magnitude estimates based on homogeneous nucleation. We emphasize that this new effect appears for the (today) most probable range of cosmological and QCD parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/hep-ph0011359_arXiv.txt": {
        "abstract": "\\noindent We compare predictions for the spin--independent contribution to the neutralino--proton scattering cross section $\\sig$ and for the anomalous magnetic dipole moment of the muon, $a_\\mu = (g_\\mu-2)/2$, in models with gravity--mediated supersymmetry breaking. We nearly always find a positive correlation between these two measurables, i.e. scenarios with larger $a_\\mu$ also tend to have larger $\\sig$, but the detailed prediction differs greatly between models. In particular, we find that for the popular mSUGRA scenario with universal soft breaking masses at the scale of Grand Unification, measurements of $a_\\mu$ currently seem more promising. On the other hand, if scalar soft breaking masses at the GUT scale receive sizable contributions from $SO(10)$ D--terms, one often finds scenarios with large $\\sig$ but $a_\\mu$ below the currently foreseen sensitivity. A string--inspired model with non--universal scalar spectrum at the GUT scale falls between these two cases. ",
        "introduction": "The Minimal Supersymmetric Standard Model (MSSM) \\cite{SUSY} is one of the best motivated extensions of the Standard Model. It offers a natural solution of the hierarchy problem \\cite{witten} as well as amazing gauge coupling unification \\cite{amaldi}. Since naturalness requires that at least some superparticles have masses at or below the TeV scale, supersymmetric theories generally predict a rich phenomenology at future colliders \\cite{colrev}. As extra ``bonus'', the simplest version of the MSSM, where $R-$parity is conserved, also contains a new stable particle (the lightest supersymmetric particle, LSP); in most cases this is the lightest neutralino, which often makes a good Dark Matter candidate \\cite{jungman}.  Unfortunately, direct searches for superparticles at high energy colliders so far yielded null results. Although it is quite possible that direct evidence for supersymmetry will be found at the upcoming ``run 2'' of the Tevatron collider, a decisive test of weak scale supersymmetry may only be possible at the LHC, which will not commence operations for another five or six years. In the meantime it is important to explore slightly less direct methods that might either give evidence for supersymmetry, or else constrain the (usually quite large) parameter space of supersymmetric models. Here we study two kinds of experiments which are expected to yield results with greatly improved sensitivity in the next couple of years: direct searches for ambient Dark Matter\\footnote{Many physicists would accept direct detection of weakly interacting massive particles (WIMPs) as (component of) the dark halo of our galaxy as fairly direct evidence for weak--scale supersymmetry. However, given that direct detection experiments can basically only measure the mass and scattering cross section of the WIMP, such a detection, while undoubtedly of the greatest importance, may not convince everybody that supersymmetry has been found.}, and the measurement of the anomalous dipole moment of the muon.  Current direct Dark Matter search experiments are sensitive to scenarios with $\\tilde{\\chi}_1^0 p$ scattering cross section $\\sig$ of order $10^{-6}$ pb or more. In fact, the DAMA experiment even claims evidence for WIMP Dark Matter, based on the annual modulation of their counting rate \\cite{DAMA}; however, the CDMS experiment, which has comparable sensitivity, sees no signal \\cite{CDMS1}. Within the next two or three years we can expect results from the upgraded CDMS experiment (at the deep underground Soudan site) \\cite{CDMS2} as well as from the CRESST experiment in the Gran Sasso laboratory \\cite{cresst}. The sensitivity of these experiments should be about two orders of magnitude better than that of current experiments, so that scenarios with $\\sig$ exceeding $10^{-8}$ pb might be testable.  SUSY models can also be tested using precision measurements at low energy experiments. Through their loop effects, sparticles contribute to low energy physics and these effects may become significant if the masses of sparticles are not too large. The supersymmetric contribution to the muon magnetic dipole moment (MDM) $a_\\mu$ = ${1 \\over 2} (g-2)_\\mu$ is one of the most robust probes. Unlike supersymmetric contributions to rare decays or CP violation in the $B-$system, which are currently being studied at the $B-$factories BaBar and BELLE, the supersymmetric contribution to $a_\\mu$ is not very sensitive to the way the supersymmetric flavor and CP problems are solved.  At present, the muon MDM is measured to be $a_\\mu^{exp} = (116~592~05 \\pm 45) \\times 10^{-10}~$\\cite{carey} and is consistent with the Standard Model. The ongoing analysis of data from the Brookhaven experiment E821 is expected to reduce the uncertainty further and the ultimate goal of the experiment is $\\Delta a_\\mu \\sim 4 \\times 10^{-10}~$\\cite{e821}.\\footnote{In order to fully exploit such a small experimental error it is also necessary to reduce the hadronic uncertainty of the SM prediction, which currently stands at $\\sim 7 \\times 10^{-10}$ \\cite{carey,kino}. Here the analysis of data that is being collected at the low--energy $e^+e^-$ colliders in Frascati, Novosibirsk and Beijing, as well as from semi--leptonic decays of the $\\tau$ lepton, will be crucial.} Supersymmetric contributions can easily exceed this sensitivity without violating any other constraints on the supersymmetric particle spectrum (from direct searches or other loop effects) \\cite{amugen}.  These supersymmetric contributions to $a_\\mu$ stem from smuon--neutralino and sneutrino--chargino loops. Their size thus depends on the parameters that appear in the chargino and neutralino mass matrix, as well as on the masses of the sleptons. Of course, the parameters of the neutralino mass matrix also determine the mass and composition of the lightest neutralino, which we assume to form the Dark Matter in the Universe. Moreover, at least for the theoretically most appealing case of a bino--like LSP, the LSP relic density is essentially determined by the mass of the LSP as well as the masses of the $SU(2)$ singlet sleptons \\cite{dn3}. The predictions for the relic density of bino--like LSPs and for the supersymmetric contribution to $a_\\mu$ are thus related.  There also exist relations between the SUSY contribution to $a_\\mu$ and the predicted LSP--nucleon scattering cross section. As already mentioned, both quantities depend on the parameters of the neutralino mass matrix. An even more direct connection comes from the common dependence on the ratio \\tanb\\ of the vevs of the two neutral Higgs fields of the MSSM, since both $a_\\mu$ and $\\sig$ increase with increasing \\tanb. The operator in the effective Lagrangian that gives rise to the anomalous magnetic moment of the muon couples left-- and right--handed muons, i.e. violates chirality. In the MSSM (as well as in the SM) all chirality violation in the (s)muon sector is proportional to the Yukawa coupling of the muon. This coupling in turn is proportional to $1/\\cos \\! \\beta$, which scales like \\tanb\\ if $\\tan^2\\beta \\gg 1$. The leading contributions to the spin--independent (coherent) contribution to $\\sig$ also involve violation of chirality \\cite{dn5}, but this time in the (s)quark sector.\\footnote{There are also squark exchange contributions where chirality violation is provided by the LSP mass, but these contributions are suppressed by an extra power of $(m^2_{\\tilde q} - m^2_{\\tilde \\chi})^{-1}$ compared to the leading squark exchange contribution, and can thus usually be ignored in the type of model we are considering \\cite{dn5}.} These contributions either come from the exchange of CP--even Higgs bosons, or from squark exchange. In both cases chirality violation ultimately comes from the quark Yukawa coupling, which, in case of down--type quarks, has the same \\tanb\\ dependence as the Yukawa coupling of the muon. Note that searches for Higgs bosons at LEP already imply $\\tanb \\gsim 2$ in the MSSM\\footnote{Higgs searches also permit $\\tanb \\lsim 0.6$, but this is strongly disfavored on theoretical grounds, since the top Yukawa coupling would have a Landau pole quite close to the weak scale.}; this implies that the Yukawa couplings of up--type quarks, which scale $\\propto 1/\\sin\\beta$, must be close to their SM--values, and essentially independent of \\tanb. The upshot of this discussion is that both the SUSY contribution to $a_\\mu$ \\cite{amugen,moroi} and $\\sig$ \\cite{sigold,dn5,signew} can be quite large if $\\tanb \\gg 1$.  In this paper, we investigate these connections between the SUSY contribution $\\Delta a_\\mu$ to the magnetic dipole moment of the muon and the spin--independent neutralino--proton cross section quantitatively. We study the minimal supergravity (mSUGRA) model as well as some other SUSY models where the assumption of strict scalar mass universality at the GUT scale is relaxed.  In Sec. 2, we discuss $\\Delta a_\\mu$ and $\\sig$ in the mSUGRA scenario.  These observables are below the sensitivity of near--future experiments in the small $\\tan\\beta$ region, partly because the experimental Higgs mass bounds then tightly restrict the allowed SUSY parameter space. However, they can become significant, and could be detectable, if \\tanb\\ is large. In Sec. 3 we relax the assumption of universal scalar masses. We find that the correlation between $\\Delta a_\\mu$ and $\\sig$ quite sensitively depends on the boundary condition at the GUT scale. Sec. 4 is devoted to conclusions. ",
        "conclusions": "In this paper, we investigated the SUSY contribution to the muon MDM $a_\\mu$, and to the spin--independent neutralino--proton cross section $\\sig$, in several SUSY models with either universal or non--universal soft scalar masses at the GUT scale. Both these observables can become significant if \\tanb\\ is large, but will be difficult to measure in the small \\tanb\\ region. Both quantities are sensitive to chirality violation in the matter (s)fermion sector, which in the MSSM is enhanced for large \\tanb. Conversely, if \\tanb\\ is small the experimental Higgs mass bounds tightly restrict the allowed SUSY parameter space, forcing the SUSY breaking scale to be quite high. While these general statements are fairly model--independent as long as one sticks to the field content of the MSSM, quantitative predictions do depend significantly on details of the spectrum of superparticles, in particular on the implementation of radiative gauge symmetry breaking.  We considered three different models: mSUGRA, a superstring scenario where SUSY is broken by the $F-$terms of the dilaton and the moduli superfields, and an $SO(10)$ GUT model. The superstring model in principle has a very large parameter space, since it allows for non--universal masses for matter sfermions and Higgs bosons, subject only to the sum rule (\\ref{string-sum}). Here we explored two orthogonal directions in this parameter space, where universality is violated either only in the Higgs sector or only in the slepton sector. We found that allowing the Higgs masses to differ from the sfermion masses at the GUT scale can vary the predicted value of $\\sig$ by a factor of a few in either direction, compared to the mSUGRA prediction for the same parameters. Given the uncertainty in the prediction of $\\sig$ due to unknown hadronic matrix elements, and the uncertainty in the predicted LSP--nucleus scattering rate due to the uncertainty of the ambient LSP flux, such a variation is barely significant. The prediction of $\\Delta a_\\mu$ is almost insensitive to this variation of the Higgs soft breaking masses. Conversely, if we maintain the universality of the soft breaking contributions to squark and Higgs masses but allow the masses of $SU(2)$ doublet and singlet sleptons to differ, we can change the prediction for $\\Delta a_\\mu$ by a factor of $\\sim 2$ around the mSUGRA value, while keeping $\\sig$ constant; the predicted LSP relic density, and hence the cosmologically preferred region of parameter space, also depends on this ratio of slepton masses.  Generally speaking, in mSUGRA as well as its stringy variant the measurement of the anomalous dipole moment of the muon seems to hold more promise in the near future than searches for LSP Dark Matter; this is largely due to the most recent severe constraints on the Higgs sector from LEP experiments.  We should caution the reader that this conclusion crucially depends on our choice to only study scenarios where both squarks and gluinos can be produced at the LHC. Given current slepton and chargino search limits, this implies that sfermion and gaugino masses are very roughly of the same order of magnitude. If one allows scalar masses $m \\gsim 1$ TeV and $m^2 \\gg M^2$, one can find solutions in mSUGRA where the LSP has a large, or even dominant, higgsino component \\cite{focus}. Since in this region of parameter space all sfermions are very heavy, the supersymmetric contribution to $a_\\mu$ is totally negligible, but $\\sig$ can be very large \\cite{focus2}.  Much more dramatic deviations from mSUGRA predictions are possible if one introduces $SO(10)$ $D-$term contributions to scalar masses. In this model the maximal allowed value of $\\sig$ for given \\tanb\\ can exceed the mSUGRA prediction by more than a factor of 30, if the $D-$term contribution to the mass of the Higgs boson that couples to the top quark is positive. For the same sign of the $D-$terms the supersymmetric contribution to $a_\\mu$ is reduced by a factor of a few, since the change of the slepton masses increases the importance of neutralino loop diagrams compared to chargino loops, leading to a strong cancellation in the total result. In this kind of model direct Dark Matter searches therefore appear to be more promising in the next few years than measurements of $a_\\mu$.  A striking result of our analysis is that even if we restrict squark and gluino masses to lie within discovery range of the LHC, and in addition fix the value of \\tanb, the prediction for $\\sig$ still varies by almost three orders of magnitude within the range of models we studied; in contrast, the prediction for $\\Delta a_\\mu$ ``only'' varies by about one order of magnitude; see Fig.~3. One reason for this difference is that $\\sig$ measures squared Feynman amplitudes, while $a_\\mu$ is directly proportional to a (loop) amplitude. In addition, we saw that $\\Delta a_\\mu$ does not depend on $\\mu$ very sensitively, since various effects tend to cancel. In contrast, the sensitivity of $\\sig$ to $\\mu$ is often enhanced by cancellations between diagrams with different $\\mu-$dependence. A similar remark holds for the \\tanb\\ dependence. Fig.~1 shows that it is ``only'' linear for $\\Delta a_\\mu$, but much stronger for $\\sig$. In this case the strong sensitivity of $\\sig$ is partly due to the dependence of the mass of the heavy CP--even Higgs boson on \\tanb, which is a direct consequence of radiative gauge symmetry breaking. Given the theoretical uncertainties involved in predicting Dark Matter detection rates, this strong parameter dependence of $\\sig$ is in some sense fortunate; it means that even an order--of--magnitude determination of $\\sig$ will allow us to make significant statements about supersymmetric parameters. Given its comparatively mild parameter dependence, it is also fortunate that the prediction of $a_\\mu$ is much cleaner, if we assume that the hadronic uncertainty can be reduced to the level of the expected experimental uncertainty, which seems feasible.  Our overall conclusion is that measurements of $a_\\mu$ and/or $\\sig$ might well yield a positive signal in the next few years. These measurements can help to distinguish between different supersymmetric models. Since well motivated models exist where one of these quantities is sizable while the other is small, it is important to measure, or at least constrain, {\\em both} of them as precisely as possible. On the other hand, we also saw that two models with very different scalar spectrum (the $SO(10)$ model and mSUGRA with $m^2 \\gg M^2$) can lead to very similar predictions for both $a_\\mu$ and $\\sig$. This once again illustrates the truism that indirect measurements can provide crucial information, but they can never completely replace direct searches for new particles.   \\subsection*{Acknowledgements} The work of M.D. was supported in part by the ``Sonderforschungsbereich 375--95 f\\\"ur Astro--Teilchenphysik'' der Deutschen Forschungsgemeinschaft. M.M.N. was supported in part by a Grant-in-Aid for Scientific Research from the Ministry of Education (12047217)."
    },
    "0011/astro-ph0011508_arXiv.txt": {
        "abstract": "As is well-known, the requirement that gamma ray bursts be optically thin to high energy photons yields a lower limit on the Lorentz factor ($\\gamma$) of the expansion.  In this paper, we provide a simple derivation of the lower limit on $\\gamma$ due to the annihilation of photon pairs, and correct the errors in some of the previous calculations of this limit. We also derive a second limit on $\\gamma$ due to scattering of photons by pair-created electrons and positrons.  For some bursts, this limit is the more stringent. In addition, we show that a third limit on $\\gamma$, which is obtained by considering the scattering of photons by electrons which accompany baryons, is nearly always less important than the second limit. Finally, we evaluate these limits for a number of bursts. ",
        "introduction": "\\label{sec:method1}   It is assumed that the photon spectrum at high energies is a truncated power law. This can be regarded as an approximation to the high energy portion of the popular \\cite{bmf+93} parameterization.  Thus, the spectrum (i.e. the observed number of photons per unit time per unit area per unit energy) is given by $fe^{-\\alpha}  , \\  e_{\\rm min}<e<e_{\\rm max}$. Here $e$ is the observed photon energy, and, unless explicitly stated otherwise, all energies throughout this paper refer to their values in the observer's reference frame. The factor $f$ is a normalization factor. The exponent $\\alpha$ is typically between 2 and 3; a value greater than 2 implies that most of the energy is in low-energy photons. The energy $e_{\\rm max}$ is the highest energy for which photons were observed. Very often, it is the  energy above which the photon flux is too small to be detected \\citep{cds98}. The spectrum  turns over below $e_{\\rm min}$, which is typically around 0.1-1 MeV, although the exact value will not be needed in this paper.  Of particular interest in the calculations below will be the total number of emitted photons which have energies greater than $e$, if the burst emits for an observed time $\\delta T$ (where $\\delta T$ is set by the variability timescale as seen in BATSE light curves, for example). This number, $N_{>e}$, is obtained by integrating the spectrum with respect to $e$, and then multiplying by $\\delta T$ and $4\\pi d^2$ (where $d$ is the distance to the GRB), yielding \\begin{equation} N_{>e}=4\\pi d^2 f\\delta T e^{-\\alpha+1}/(\\alpha-1) \\ . \\label{eq:ngte} \\end{equation} Throughout this paper, we assume spherical symmetry (at least within the beaming angle $1/\\gamma$). In addition, for clarity, we neglect corrections due to cosmological redshift in the text. However, these corrections will be included in table~(1) which summarizes the main results.  Now, in the frame of the emitting material, where the photons are assumed to be roughly isotropic, a photon with energy $e'$ can annihilate a second photon with energy larger than $(m_e c^2)^2/e'$, yielding an electron-positron pair, where $m_e$ is the electron mass. When the energy of the second photon is around $(m_e c^2)^2/e'$, then the cross-section for this process is approximately the Thomson cross-section, $\\sigma_T$. The cross-section falls off as a power-law of the annihilating photon energy if its energy is significantly larger than this value, and it is zero when its energy is below this value. The combination of these effects leads to an averaged cross section of about $\\frac{11}{180}\\sigma_T$, assuming an $\\alpha=2$ spectrum \\citep{sve87}.  If the emitting material is moving towards the observer with a Lorentz factor $\\gamma$, then the photons are blueshifted by $\\gamma$. Thus, a photon with detected energy $e=\\gamma e'$ can only annihilate photons whose detected energies are larger than $(\\gamma m_e c^2)^2/e$. Since most of the photons are at small energies, the photon with the highest energy will be most susceptible to annihilation by other photons. Thus, a lower limit on the Lorentz factor can be obtained by requiring that the photon with energy $e_{\\rm max}$ will have optical depth smaller than unity.  To calculate this optical depth, we note that when the expansion is spherically symmetric (at least within the beaming angle $1/\\gamma$), the burst expands to a radial distance of about $\\gamma^2 c \\delta T$ \\citep{pir99}. This is true (within numerical factors of order unity) independent of the details of either the expansion scenario or of whether the emission is from internal or external shocks.  Now, if $e_{\\rm max,an}$ is defined as the energy of the photon which annihilates $e_{\\rm max}$, i.e. $e_{\\rm max,an}\\equiv (\\gamma m_e c^2)^2/e_{\\rm max}$, then there are $N_{>e_{\\rm max,an}}$ photons which can annihilate the $e_{\\rm max}$ photon. Thus, the optical depth is \\begin{equation} \\tau={{11\\over 180}\\sigma_TN_{>e_{\\rm max,an}} \\over 4\\pi(\\gamma^2c\\delta T)^2 }\\ . \\label{eq:tau} \\end{equation}  Finally,  inserting the expression for $N_{>e}$ (equation~\\ref{eq:ngte}) into the above equation yields \\begin{equation} \\tau= \\hat{\\tau} \\Big( {e_{\\rm max}\\over m_e c^2} \\Big)^{\\alpha-1} \\gamma^{-2\\alpha-2} \\ , \\label{eq:tau1} \\end{equation} where \\begin{equation} \\hat{\\tau}\\equiv {{11\\over 180}\\sigma_T d^2  (m_e c^2)^{-\\alpha+1}f\\over c^2 \\delta T (\\alpha-1)} \\label{eq:tauhat} \\end{equation} is a dimensionless quantity which will appear again below. Physically, it is the optical depth for a photon with energy $m_ec^2$ in a mildly relativistic fireball (i.e. with expansion speed $v\\approx c$,  but with $\\gamma$ not much larger than unity). It is evaluated with observationally plausible values in equation~(\\ref{eq:numtauhat}).  The requirement that $\\tau<1$ leads to the limit \\begin{equation} \\gamma>\\hat{\\tau}^{1\\over 2\\alpha+2}\\Big( {e_{\\rm max}\\over m_ec^2}\\Big) ^{\\alpha-1\\over 2\\alpha+2} \\ . \\label{eq:gamma1} \\end{equation}   The interpretation of the above calculation can be facilitated by defining a characteristic photon energy, $e_{\\rm thick}$. This energy is defined as follows: the number of photons  with energy greater than $e_{\\rm thick}$ is such that, if each of these photons were assigned a cross-section equal to $11/180$ of the Thomson cross-section, then they would be optically thick, i.e., \\begin{equation} 1\\equiv{ {11\\over 180} \\sigma_TN_{>e_{\\rm thick}} \\over 4\\pi(\\gamma^2c\\delta T)^2} \\ , \\ {\\rm so\\  that} \\ \\ {e_{\\rm thick}\\over m_e c^2}=\\hat{\\tau}^{1\\over \\alpha-1} \\gamma^{-{4\\over \\alpha-1}} \\ . \\label{eq:ethick} \\end{equation} Since the number of photons decreases with increasing energy, the requirement that the optical depth of the photon with energy $e_{\\rm max}$ be less than unity is simply the requirement that $e_{\\rm max,an}>e_{\\rm thick}$---from which equation~(\\ref{eq:gamma1}) follows directly.  Table (1) summarizes the main results of this paper. In this table, the corrections due to cosmological redshift ($z$) have been added. These corrections may be obtained by considering equation~(\\ref{eq:tau}).  There are three redshift effects. (i) In calculating $N_{>e_{\\rm max,an}}$, one needs to convert from the number of observed photons per unit area to the total number of photons at the source: instead of multiplying by $4\\pi d^2$, one should multiply by the appropriate surface area of a sphere which is centered at the burst, i.e. by $4\\pi [d_L/(1+z)]^2$, where $d_L$ is the luminosity distance. (ii) Due to cosmological time dilation, $\\delta T\\rightarrow \\delta T/(1+z)$ in the denominator of equation~(\\ref{eq:tau}). Note that when (i) and (ii) are combined, the explicit factors of $1+z$ cancel, although there is still a redshift dependence through the luminosity distance. (iii) Due to the redshift of photon energies, the photon which annihilates the $e_{\\rm max}$ photon need only have an energy greater than $((\\gamma/(1+z) m_e c^2)^2/e_{\\rm max}$ in order to exceed the threshold for pair creation, i.e. $e_{\\rm max,an}\\rightarrow e_{\\rm max,an}/(1+z)^2$. When this is inserted this into $N_{>e_{\\rm max,an}}$ in equation~(\\ref{eq:tau}), and combined with (i) and (ii), the result is that $\\tau\\propto d_L^2(1+z)^{2\\alpha-2}$.  \\subsection{Comparison With Other Work}  The physical mechanism presented in this section has been used in a number of previous papers to place a lower limit on the Lorentz factor.  However, in most of these papers, we found that the dependence of the optical depth on the Lorentz factor was incorrect. This incorrect dependence results in a limit on $\\gamma$ which is too large by a factor of, typically, 2 or 3. The correct dependence is given in equation~(\\ref{eq:tau1}), i.e., $\\tau\\propto\\gamma^{-2\\alpha-2}$.  Both \\cite{feh93} and \\cite{wl95} considered the case when $\\alpha=2$, and both obtained $\\tau\\propto\\gamma^{-4}$ instead of the above scaling $\\tau\\propto\\gamma^{-6}$ when $\\alpha=2$. (See figure~(2) in the former reference and equation~(12) in the latter reference.) Both of these papers used an incorrect expression for the dependence of the optical depth on the cross-section.  In equation~(5) of \\cite{feh93}, and in equation~(4) of \\cite{wl95}, there should be an extra factor of $1-\\cos\\theta$ multiplying the cross-section, where $\\theta$ is the angle between annihilating photons.  This extra factor is in addition to the dependence of the cross-section on $\\theta$. It accounts for the fact that, if two photons are traveling nearly parallel to each other, then they are unlikely to interact---in the limit that they are exactly parallel ($\\theta=0$), they will never interact.  Since, approximately, $1-\\cos\\theta=\\gamma^{-2}$, this accounts for the difference between these authors' results and ours.   \\cite{bh97} presented an expression for the optical depth which is proportional to $\\gamma^{-2\\alpha}/R_0$, where $R_0$ is the radius of the burst (their equation~(41), for example). This is in agreement with our expressions. However, they then claimed that $R_0$ is proportional to ``one or two powers'' of $\\gamma$. They finally concluded, in their section~3.3, that $R_0\\simeq\\gamma c \\delta T$ is more appropriate, yielding $\\tau\\propto \\gamma^{-2\\alpha-1}$, and used that for their numerical results. However, as long as the burst is spherically symmetric within the angle $1/\\gamma$, the only possibility is $R_0\\simeq\\gamma^2 c \\delta T$. While the numerical coefficient in this equation is uncertain, the exponent on $\\gamma$ is not. The relevant quantity is the number of photons per unit area at the source. One could also use $R_0\\simeq\\gamma c \\delta T$, which is the transverse size seen by a given observer, in the denominator of equation~(\\ref{eq:tau}), but then only the photons emerging from the transverse area should be used in the numerator.  The number of these photons is correspondingly reduced by $1/\\gamma^2$, and our equation~(\\ref{eq:tau}) remains valid. In addition, we note that the dependences on cosmological redshift are incorrect in \\cite{bh97}.  The correct dependence for the optical depth is $\\tau\\propto (1+z)^{2\\alpha-2}$, assuming that it is the luminosity distance which is used. The resulting dependence on redshift of the lower limit on $\\gamma$ is as listed in table~(1).  Our calculation is similar in spirit to that presented in the review article of \\cite{pir99} (his equation~(10)). However, when calculating the number of photons, he used the total number of photons emitted throughout the burst, whereas it is more appropriate to consider only the number of photons emitted during the variability timescale $\\delta T$. After correcting for this, his estimate of the optical depth is similar to ours (assuming $\\alpha$ in his equation~(10) is our $\\alpha-1$, rather than as defined in his equation~(2)).  Finally,  \\cite{kp91} calculated the optical depth, but in a scenario which is different than ours. Specifically, they assumed that the emitting material consists of rigid blobs which move with a bulk Lorentz factor $\\gamma$, such as would be the case for a relativistically moving star.  Therefore, their limit depends only on the luminosity and not on the timescale of variability. They obtained $\\tau\\propto\\gamma^{-2\\alpha-1}$, which is correct for their scenario, but which differs from our result by one power of $\\gamma$.  While the aforementioned papers presented more detailed scenarios for the expansion of the burst than presented here, differing scenarios would change our parameter $\\hat{\\tau}$ by factors which are of order unity.  However, the dependence of the minimum $\\gamma$ on $\\hat{\\tau}$ is very weak---$\\gamma$ is proportional to  $\\hat{\\tau}^{1/6}$ when $\\alpha=2$---and thus these factors of order unity have little effect on the limit on $\\gamma$. We prefer to leave the expansion scenario vague, largely because there are too many alternatives, each of which would produce a different factor of order unity in $\\hat{\\tau}$, and none of which would significantly affect the minimum $\\gamma$.  In addition, when we compared our result with the more detailed model of  \\cite{bh97} (but with their $\\gamma$-dependence corrected), we found that our lower limit on $\\gamma$ exceeded theirs by less than around 30\\% for typical values of $\\alpha$.  \\subsection{Discussion of Two Assumptions} \\label{sec:dota}  In the above calculation of the limit on $\\gamma$, two implicit assumptions have been made.  First, it was assumed that $e_{\\rm max,an}>e_{\\rm min}$, i.e. that the photon which annihilates the $e_{\\rm max}$ photon does not have an energy which is smaller than the low-energy break in the photon spectrum. In terms of observed quantities, this condition can be written as $e_{\\rm max}<m_ec^2\\hat{\\tau}^{1/2}$ (assuming that $e_{\\rm min}<m_ec^2$). It is easily satisfied for all of the bursts which we have considered (except for the burst from which TeV photons may have been detected;  see section~(\\ref{sec:obsres}), below). If it is not satisfied, then the limit on $\\gamma$ depends on the behaviour of the photon spectrum below $e_{\\rm min}$.  In particular, if the number of photons is dominated by the high end of this part of the spectrum (i.e. by photons whose energies are around $e_{\\rm min}$), then the condition to be optically thin is $e_{\\rm thick}<e_{\\rm min}$; this leads to $\\gamma> \\hat{\\tau}^{1\\over 4}/ (e_{\\rm min}/m_e c^2)^{\\alpha-1\\over 4}$. Conversely, if the number of photons is dominated by the low end of this part of the spectrum, then the limit which has been calculated in this section in equation~(\\ref{eq:gamma1}) is applicable as long as it is this part of the spectrum which is used to calculate $\\hat{\\tau}$;  i.e. in equation~(\\ref{eq:tauhat}), the quantities $f$ and $\\alpha$ should be defined by requiring that the spectrum below $e_{\\rm min}$ is $fe^{-\\alpha}$.  Second, it was assumed that the photon with energy $e_{\\rm max}$ can annihilate a second photon whose energy is equal to its own (i.e. $e_{\\rm max}>m_ec^2\\gamma$), and thus that photons which annihilate the $e_{\\rm max}$ photon have a minimum energy which is less that $e_{\\rm max}$.  If, conversely, the lower limit on $\\gamma$ obtained in equation (\\ref{eq:gamma1}) does not satisfy this constraint (i.e. if $\\gamma>e_{\\rm max}/m_ec^2$), then the $e_{\\rm max}$ photon can only annihilate photons with energies larger than itself. Since we do not observe these photons, we can only speculate about their existence.  We consider two alternatives.  A first alternative is that the GRB does not produce any photons whose energies would be observed to be larger than $e_{\\rm max}$ (independent of any optical depth considerations). This alternative implies that the limit on the Lorentz factor is \\begin{equation} \\gamma>e_{\\rm max}/m_ec^2 \\ . \\end{equation} However, it seems unlikely that this limit  would be valid for most bursts: since $e_{\\rm max}$ is often the energy above which the extrapolated photon flux would be too small to be observed, it would be overly coincidental if $e_{\\rm max}$ were also the energy above which the intrinsic photon spectrum (i.e. before optical depth considerations) were cut off.  A second alternative is that this intrinsic spectrum has no cutoff at very high energies. We consider this to be the more realistic of the two alternatives. In this case, the limit calculated previously in equation~(\\ref{eq:gamma1}) would be valid. However, if this second alternative is true, then there is a more stringent bound on $\\gamma$ which will be discussed in the following section. ",
        "conclusions": "We derived three limits on the Lorentz factor $\\gamma$, based on the requirement that GRB's be optically thin. Our order of magnitude calculations are unaffected by the details of the scenario and our lower limits apply to internal as well as external shocks, assuming spherical symmetry within an opening angle of $1/\\gamma$. Limit~A was obtained by considering the annihilation of pairs of photons; limit~B was obtained by considering the scattering of photons by pair-created e$^\\pm$; limit~C was obtained by considering the scattering of photons by electrons which are associated with baryons.  It was shown that, as long as the intrinsic photon spectrum (i.e. before optical depth effects) can be extrapolated to very high energies, limit~C is nearly always less important than limit~B. Table~(1) summarizes the results for limits~A and~B.  We evaluated limits~A and~B for a few selected bursts (see table~2).  Our numerical results for limit~A are not very different from the results of \\cite{wl95} and \\cite{bh97}. While correcting their dependence of optical depth on Lorentz factor reduced the lower limit on $\\gamma$ by a factor of a few, our use of larger distances (because of recent redshift measurements) partially compensated for this.  With the anticipated launch of the Gamma Ray Large Area Space Telescope (GLAST) in 2005, the limits presented in this paper may be improved. Since cut-offs in the photon spectrum at high energies have usually not been observed, the values of $e_{\\rm max}$ which we have used in this paper are lower limits, set by instrumental sensitivity. GLAST will be able to detect photons up to 300 GeV, with sensitivity much better than that of EGRET (http://glast.stanford.edu/mission.html). If GLAST does not observe a cut-off, then compared with our assumed value for the maximum observed energy photon from bursts with measured resdshifts (i.e. $e_{\\rm max}=m_ec^2$), $e_{\\rm max}$ would be increased by a factor of nearly $10^6$ .  Since our limit A is proportional to $e_{\\rm max}^{1/6}$ (when $\\alpha=2$), our limits could be increased by a factor of around 10, yielding lower limits on Lorentz factors of nearly 1000."
    },
    "0011/astro-ph0011022_arXiv.txt": {
        "abstract": "We use optical rotation curves to investigate the $R$-band Tully-Fisher properties of a sample of 90 spiral galaxies in close pairs, covering a range of luminosities, morphological types, and degrees of tidal distortion. The galaxies follow the Tully-Fisher relation remarkably well, with the exception of eight distinct $\\sim$3$\\sigma$ outliers. Although most of the outliers show signs of recent star formation, gasdynamical effects are probably the dominant cause of their anomalous Tully-Fisher properties.  Four outliers with small emission line widths have very centrally concentrated line emission and truncated rotation curves; the central emission indicates recent gas infall after a close galaxy-galaxy pass. These four galaxies may be local counterparts to compact, blue galaxies at intermediate redshift.  The remaining galaxies have a negligible offset from the reference Tully-Fisher relation, but a shallower slope (2.6$\\sigma$ significance) and a 25\\% larger scatter. We argue that triggered star formation is a significant contributor to the slope difference. We characterize the non-outlier sample with measures of distortion and star formation to search for third parameter dependence in the residuals of the TF relation. Severe kinematic distortion is the only significant predictor of TF residuals; this distortion is not, however, responsible for the slope difference from the reference distribution.  Because the outliers are easily removed by sigma clipping, we conclude that even in the presence of some tidal distortion, detection of moderate ($\\gtrsim 0.5$ magnitudes in rest-frame $R$) luminosity evolution should be possible with high-redshift samples the size of this 90-galaxy study. The slope of the TF relation, although difficult to measure, is as fundamental for quantifying luminosity evolution as the zero-point offset. ",
        "introduction": "Tully-Fisher (1977; TF hereafter) luminosity-linewidth studies frequently exploit the TF relation as a secondary distance indicator. As the fundamental scaling relation for spiral galaxies, the TF relation also provides constraints on galaxy formation; it is deeply connected to the processes by which disk galaxies form (e.g., Cole et al. 1994; Eisenstein \\& Loeb 1996; Navarro \\& Steinmetz 2000; Mo \\& Mao 2000). By allowing comparisons of galaxies at high $z$ to {\\it physically similar} systems at the present day, the TF relation for spirals and the fundamental plane for spheroid-dominated systems provide crucial anchors for measuring luminosity evolution directly.  Both the TF and the fundamental plane relations are indispensable for establishing a complete picture of galaxy evolution. Kinematic estimates are frequently necessary for ascertaining the masses of galaxies. For example, the faint, compact, blue galaxies may be small, starbursting galaxies that will eventually fade or larger galaxies with only moderate amounts of star formation (Koo et al. 1994, 1995; Kobulnicky \\& Zaritsky 1999). Only kinematic linewidths that truly reflect the masses of these galaxies would resolve the degeneracy. To date, TF studies of galaxies at large redshift with resolved rotation curves contain few galaxies and yield discrepant results; Simard \\& Pritchet (1998) find an offset of 1.5~--~2  magnitudes, while Vogt et al. (1996; 1997) find evidence for only a moderate offset ($\\lesssim 0.4$ magnitudes in $B$). With the advent of 8-m class telescopes, the size of these studies will increase and an understanding of the underlying physics will be crucial. Until now, the primary goal of large TF studies has been construction of the tightest relation possible for use as a distance indicator; thus most local TF studies are limited in morphology, excluding galaxies with signs of interaction and/or tidal distortion (e.g., Rubin et al. 1985; Pierce \\& Tully 1992; Willick et al. 1995; Willick et al. 1996; Bureau, Mould, \\& Staveley-Smith 1996; Giovanelli et al. 1997a; Giovanelli et al. 1997b; Giovanelli et al. 1997c; Courteau 1997; Dale et al. 1997; Haynes et al. 1999a; Haynes et al. 1999b; Dale et al. 1999; Tully \\& Pierce 2000).  With application to high redshift in mind, we explore the effects of loosening the morphological and environmental constraints that the largest studies apply. Pruning of high-$z$ TF samples is more difficult because interactions are more frequent and the signatures of interaction are not so readily observable.  Low surface brightness companions, minor merger remnants embedded in larger galaxies, and faint tidal features can fall below the detection limits. Seeing effects obscure both morphological and kinematic distortion. The desired measurement of luminosity evolution may become confused by the effects of tidal distortion on the measurements of the galaxy parameters. Thus, high-redshift studies must rely on knowledge of the statistical effects of interactions gleaned from low-redshift studies.  The TF relation has not yet been extensively studied as a means of detecting evolution. We consider the slope of the TF relation and the outliers to the distribution, in addition to the zero-point offset, as fundamental tools for characterizing luminosity evolution and morphological evolution.  We explore the TF properties of 90 galaxies in the Barton, Geller, \\& Kenyon (2000a; BGK hereafter) sample of galaxies in close pairs and n-tuples. Because BGK find the distinct signature of increased star formation due to interactions, the BGK sample provides an ideal testing ground for the use of the TF relation to detect moderate luminosity evolution. We lay the groundwork for future TF studies of high-reshift galaxies by: (1) quantifying the effects of moderate distortion on the, intercept, slope and outliers of the TF relation, and (2) setting limits on the ability of pre- or non-merger interaction to initiate luminosity evolution off the TF relation at the current epoch. Thus, our study complements searches for TF deviations in low-mass or low-surface brightness galaxies (Sprayberry et al. 1995; Courteau \\& Rix 1999; O'Neil et al. 2000; McGaugh et al. 2000), extreme late-type galaxies (Matthews, van Driel, \\& Gallagher 1998), S0s (Neinstein et al. 1999), and asymmetric galaxies (Zaritsky \\& Rix 1997), and for non-linearity in the relation (Mould, Han, \\& Bothun 1989).  In Sec.~2, we describe the sample and the data reduction and analysis procedures.  We discuss the TF properties of paired galaxies and identify the outliers to the TF relation in Sec.~3.  In Sec.~4, we explore the expected effects of tidal and kinematic distortion and luminosity evolution; we interpret the outliers in light of these expectations and identify 4 outliers with similar systems at intermediate redshift, the compact, blue galaxies with narrow emission lines. In Sec.~5, we develop empirical measures of distortion and luminosity evolution; we test these measures for correlations with the TF residuals, and we examine causes of the differences between the the pair TF distribution and the reference TF distribution. We discuss the implications of our results for high-redshift evolution studies in Sec.~6 and we conclude in Sec.~7. ",
        "conclusions": "We use optical emission line rotation curves to investigate the $R$-band Tully-Fisher properties of a sample of 90 spiral galaxies in close pairs. The sample includes a range of luminosities, morphological types, and degrees of tidal distortion. With the exception of eight distinct $\\sim$3$\\sigma$ outliers, the galaxies follow the Tully-Fisher relation remarkably well, even in the presence of striking tidal distortion. Although most of the outliers show signs of strong star formation during the last $\\sim$100 -- 300 Myr, gasdynamical effects are probably the dominant cause of their anomalous Tully-Fisher properties.  Strong effects from morphological distortion and/or dissipationless kinematic distortion are rare.  Four outliers with small emission line widths have very centrally concentrated emission and truncated rotation curves. Recent gas infall after a close galaxy-galaxy pass or a minor merger can trigger this emission and can enhance the bulges of spiral galaxies. These four galaxies may be local counterparts to the compact, blue galaxies observed at intermediate redshift, including the compact narrow emission line galaxies.  The remaining galaxies have only a small zero-point offset from the Courteau (1997) TF relation for more isolated galaxies, but the pairs have a shallower slope (2.6$\\sigma$ significance) and a 25\\% larger scatter. We argue that triggered star formation is one significant contributor to the slope difference. We characterize the non-outlier sample with measures of distortion and star formation to search for third parameter dependence in the residuals of the TF relation. Severe kinematic distortion is the only significant predictor of TF residuals, although this distortion is not responsible for the slope difference.  Our results have several implications for TF studies at moderate or high redshift: \\begin{enumerate}  \\item Morphological distortion alone rarely results in strong effects on the TF properties of galaxies. Gasdynamical effects on the rotation curve are much more common.  These effects together give rise to a small set of outliers ($\\sim$10\\% of our sample).  \\item Outliers due to interaction-induced distortion are easily removed by sigma clipping. This clipping provides a method for identifying galaxies with concentrated emission in an epoch of bulge enhancement, such as the compact, narrow emission line galaxies.  \\item Luminous galaxies require prohibitively large bursts of star formation to move off the TF relation; at least locally, bursts of star formation will affect only low-mass galaxies. Therefore, the slope of the TF relation, while difficult to measure, is at least as fundamental for quantifying luminosity evolution as the zero-point offset.  \\end{enumerate}  We conclude that detection of moderate ($\\gtrsim 0.5$ magnitudes in rest-frame $R$) luminosity evolution should be possible with high-redshift samples the size of this 90-galaxy study, even in the presence of some tidal distortion."
    },
    "0011/astro-ph0011214_arXiv.txt": {
        "abstract": "At present, the ground-based astronomy of very high energy (VHE) (E$>$100~GeV) $\\gamma$-rays experiences the challenging transition phase caused by the substantial upgrading of its observational instrumentation. Recently the HEGRA collaboration has evidently demonstrated the power of arrays of imaging atmospheric \\v{C}erenkov telescopes and initiated a number of complementary projects, such as H.E.S.S., CANGAROO IV, VERITAS. The common philosophy of those projects is based on the stereoscopic observations of air showers with a number of imaging telescopes simultaneously. Such observations allow to reduce substantially the effective energy threshold of detected $\\gamma$-ray showers; to improve the angular and energy resolution for individual $\\gamma$-rays; to gain in suppression of the cosmic ray background. In my talk I have summarized the general advantages and achieved sensitivity of the stereoscopic observations using the HEGRA data taken with an array of 5 imaging telescopes having the energy threshold of 500 GeV. These results completed with the relevant simulations may be extrapolated on a firm ground towards the future H.E.S.S. (phase I) array of four 12 m imaging telescopes with the estimated energy threshold below 100 GeV. Here I discuss in short the anticipated performance of such an array. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011164_arXiv.txt": {
        "abstract": "By means of wide--field ($42\\times28$ arcmin) CCD images of the Coma cluster of galaxies and of control fields, taken at the Canada--France--Hawaii telescope with the CFH12K and UH8K cameras, we determine the luminosity function bi--variate in central brightness. We found a clear progression for a steeping and weakening of the FL going from high surface brightness galaxies ($\\mu\\sim20$ mag arcsec$^{-2}$) to galaxies of very--faint central brightnesses ($\\mu\\sim24.5$ mag arcsec$^{-2}$). Compact galaxies, usually rejected in the star/galaxy classification, are found to be a minor population in Coma. A huge population of low surface brightness galaxies red and faint are discovered, representing the largest contributor to the LF at faint magnitudes. Among them, there could be the remnant of high redshift starbursting galaxies responsible for the Butcher--Oemler effect. ",
        "introduction": "The luminosity function (LF hereafter), i.e. the number density of galaxies having a given luminosity, is critical to many observational and theoretical problems (see e.g. Binggeli, Sandage \\& Tammann 1988). From an observational point of view, the LF is the natural ``weight\" of all those quantities which need to be weighted against the relative number of objects in each luminosity bin. Furthermore, due to the role played by luminosity and surface brightness in the inclusion of objects in any observed sample (faint objects or low surface brightness galaxies are often excluded or under--represented), the knowledge of the LF and the LF bi--variate in brightness is fundamental to compute the selection function and is needed to derive the actual galaxy properties from the measured quantities (see, for example, the discussion on the field LF steepness by Sprayberry et al. 1997).  The LF buries under its foots the true problem (Sandage 1994): the LF is likely the sum of the LFs of the specific types, or of any other physically based galaxy classes. Galaxies can be classified on the bases of their central brightness, and in fact depending on their central brightness they lay in different parts of the Fundamental Plane (e.g. Bender, Burstein \\& Faber 1997), showing that this classification is not an aestheatical one but reflects some physical difference between the classes. Therefore, ``it would be of great importance to know what the luminosity function looks like when divided into classes of surface brightness\" (Kron 1994).  In the present paper we shortly present the LF bi--variate in surface brightness of a sample of galaxies in the Coma cluster of more than 4000 members and in three filters, down to the magnitude of three bright globular clusters, complete in surface brightness down to the brightness of the faintest cataloged LSB galaxies, spread up over the largest cluster area ever observed with CCDs. The interested reader may found details in Andreon \\& Cuillandre (2000). ",
        "conclusions": ""
    },
    "0011/astro-ph0011487_arXiv.txt": {
        "abstract": "{ We present X-ray observations of PSR 0540--69.3 with the XMM-Newton observatory. The spectra obtained with the Reflection Grating Spectrometer reveal, for the first time, emission from ionized species of O, Ne and Fe originating from the SNR shell. Analysis of the emission line spectrum allows us to derive estimates of the temperature, ionization timescale, abundances, location, and velocity of the emitting gas. ",
        "introduction": "SNR 0540--69.3 is one of only a few Crab-like supernova remnants (SNRs) with a shell. SNR0540--69.3 harbors a pulsar, PSR B0540--69.3, discovered in the soft X-ray band by Seward, Harnden, \\& Helfand (\\cite{seward1}). Like the Crab, it shows plerion emission from a synchrotron nebula powered by the embedded young pulsar (Chanan, Helfand \\& Reynolds \\cite{chanan}). From the period (50 ms) and spin down rate ($4.79 \\times 10^{-13} {\\rm s\\, s^{-1}}$) a a characteristic age of 1660 yr and rotational energy loss rate of $1.5{\\times}10^{38}$ergs s$^{-1}$ have been derived (Seward, Harnden, \\& Helfand \\cite{seward1}).  Early optical observations (Mathewson et al. \\cite{mathewson}) classified SNR 0540--69.3 as a young oxygen-rich SNR with an 8\\arcsec\\ diameter shell, bright in [\\ion{O}{iii}]. Mathewson et al. (\\cite{mathewson}) also found a diffuse patch of [\\ion{O}{iii}] $\\sim$30\\arcsec\\ to the west of the ring which indicated that the remnant was larger than 8\\arcsec. A {\\it ROSAT} High Resolution Imager observation by Seward \\& Harnden (1994) also revealed emission from well outside the central region, which they interpreted as from a patchy outer shell with a diameter of 55\\arcsec. This shell contributes $\\sim$20\\% to the measured flux in the {\\it ROSAT} 0.1--2 keV band. More recently Gotthelf \\& Wang (\\cite{gotthelf}) presented a high resolution {\\it Chandra} HRC observation of PSR B0540--69.3 which clearly shows emission from the outer shell.  In this letter, we report the results of observations of SNR 0540--69 acquired with the Reflection Grating Spectrometer on board the XMM-Newton Observatory (den Herder et al. \\cite{denherder}). These observations allow us for the first time to detect and identify emission lines from the remnant shell. Our spectral analysis allows us to derive values for the temperature, ionization timescale, abundances, location, and velocity of the emitting gas. ",
        "conclusions": "We have detected resolved X-ray line emission from the supernova remnant shell around PSR B0540--69.3, confirming the nature of the shell-like structure seen in previous experiments as a thermal SNR. Unfortunately, given the relatively low signal-to-noise of the observation, the fact that the object is only partially spatially resolved, and the uncertainties in the interpretation of the estimates of the various observable parameters, we can not uniquely decide what the evolutionary status of the remnant is, or what the relation of the line emitting gas is to the overall structure of the remnant."
    },
    "0011/astro-ph0011352_arXiv.txt": {
        "abstract": "The scattering by the inner accretion disk of X-ray radiation generated near the surface of a spinning neutron star in a low-mass X-ray binary (LMXB) has observable effects on the waveforms of millisecond X-ray flux oscillations produced e.g. during type-I bursts or in the millisecond pulsar SAX J1808.4--3658. We study these effects in the framework of a simplified model in which there is a single emitting spot on the stellar surface, which is visible both directly and in X-rays scattered from the disk. The main signature of scattering from a thin disk is that the pulse of scattered flux leads (if the star rotates in the same sense as the disk) or lags (in the contrary case) the primary pulse of direct emission by a quarter of a spin cycle. This is caused by Doppler boosting of radiation in the sub-relativistic Keplerian flow. The disk-scattered flux is revealed better in energy-resolved waveforms and the phase dependence of the polarized flux component. The phenomenon discussed permits direct testing of the presence of standard thin disks near the neutron stars in LMXBs and should be observable with future X-ray timing experiments having a few times better sensitivity than RXTE and also with sensitive X-ray polarimeters. ",
        "introduction": "The inner regions of the accretion disk around a weakly magnetic neutron star in a low-mass X-ray binary (LMXB) system intercept and re-emit a significant fraction of radiation generated near the surface of the neutron star. About half of the X-ray luminosity liberated via accretion originates in a boundary layer of the disk or in a spreading layer on the stellar surface. During type-I X-ray bursts, the source luminosity almost entirely originates on the stellar surface. Therefore, scattering of the central emission in the inner disk should always have an important effect on the properties of the observed X-ray signal (Lapidus \\& Sunyaev 1985).  The Rossi X-ray Timing Explorer (RXTE) detections of millisecond X-ray flux oscillations from many sources have demonstrated that the neutron stars in LMXBs are rapidly spinning, with frequencies ranging between 300 and 600~Hz (see Strohmayer 2000; Van der Klis 2000 for reviews). These oscillations are likely the result of azimuthal asymmetry in the generation of X-rays, e.g., there is a single emitting spot on the star (see Fig.~\\ref{sketch}). In future X-ray timing experiments, a wealth of information on the geometrical and physical properties of neutron stars and the inner accretion flows in LMXBs will be obtained from quality measurements of oscillation waveforms. These can be non-sinusoidal for several reasons. First, if one considers the situation (as is frequently done) in which there is {\\sl no accretion disk} near the neutron star, then only half (neglecting aberration of light) of the surface of the spinning star can be visible an any instant. Therefore, periodic eclipses of the emitting zone may take place. Second, since the spin velocity at the stellar equator is of the order of $0.1 c$ (the speed of light), the emission pulse should be asymmetric due to Doppler boosting (the spot repeatedly moves towards and away from us --- see Fig.~1; Chen \\& Shaham 1989; Weinberg et al. 2000). Both general relativistic light bending and finite spot size act to reduce oscillation amplitudes (Pechenick et al. 1983; Wood et al. 1988; Riffert \\& Meszaros et al. 1988; Weinberg et al. 2000).  \\begin{figure*} \\epsfxsize=\\hsize \\epsffile{sketch.ps} \\caption{Schematic picture of scattering of radiation emitted by a spot on the surface of a spinning neutron star from a surrounding Keplerian accretion disk. } \\label{sketch} \\end{figure*}  If we take into account the accretion disk lying near the neutron star, it will still be possible for the above phase-dependent effects to take place, but additional effects will appear at the same time. In this paper we study the {\\sl role of the accretion disk} in forming oscillation waveforms. More specifically, we show that there are clear observational effects that allow us to test the presence of a standard thin disk near neutron stars in LMXBs. First, such a disk screens half of the star from our view. Second, if the emitting zone is located above the disk (with respect to us), part of its flux is scattered by the disk in our line of sight.  In the innermost regions of the disk, matter is rotating with speeds $\\sim 0.5$~c. Therefore, scattered photons are Doppler-boosted along the direction of the disk rotation: i.e., towards us when the bright spot is irradiating the sector of the disk in which matter approaches us (which approximately corresponds to the situation presented in Fig.~\\ref{sketch}), and away from us when the opposite sector of the disk is being irradiated (half a cycle later). As a result, the phase dependence of the scattered flux component will be distinctly different from that of the direct emission of the spot. Namely, the pulse of scattered emission will {\\sl lead} the pulse of direct flux by roughly a quarter of a cycle. Furthermore, X-ray flux registered during this secondary pulse will be substantially harder and more polarized (in absolute units) than at other phases (Fig.~\\ref{profile}a). These conclusions only apply to binaries in which the neutron star rotates in the same sense as the disk. If the star rotates in the opposite sense (Fig.~\\ref{profile}b), the scattered component will {\\sl lag} the direct signal, by the same quarter of a cycle. This makes it possible to uncover counterrotating LMXBs (see Sibgatullin \\& Sunyaev 2000 on possibilities of formation of such systems).  \\begin{figure*} \\epsfxsize=\\hsize \\epsffile[0 160 590 700]{profile.ps} \\caption{(a) Spin-phase dependences of various characteristics of emission from a LMXB, taking into account the effect of disk-scattering, for the basic model ($R=\\rin=6GM/c^2$), $i=60^\\circ$, $\\theta=60^\\circ$, a spin frequency of $f=300$~Hz, a point-like emitting spot ($\\alpha=0$) and corotation of the star and disk. The dashed curves correspond to direct emission from the spot, the dash-dotted lines to the disk-scattered radiation, and the solid lines to the sum of these two components. The panels from top to bottom are for: photon flux, mean photon energy increment, polarized flux (in the same units as total flux) and polarization angle. Note that only the scattered component contributes to the polarized flux. (b) Same as (a), but for the case of star-disk counterrotation. } \\label{profile} \\end{figure*}  General relativity predicts the existence of a gap between the last stable Keplerian orbit and the stellar surface for some (soft) neutron star equations of state (see, e.g., Kluzniak \\& Wilson 1991). In geometries with a gap, part of the lower stellar hemisphere is visible, and also less flux from the upper hemisphere is scattered in the disk. However, recent calculations (Sibgatullin \\& Sunyaev 2000) show that there will be no gap if the neutron star rotates with frequency $\\sim 600$~Hz in the same sense as the disk. For the slower rotation at 300~Hz, a thin gap can exist, but its width proves to be smaller than the height of the boundary/spreading layer for the accretion rates typical of LMXBs (Inogamov \\& Sunyaev 1999; Popham \\& Sunyaev 2000). On the other hand, a large gap can arise in the case of counterrotation. We discuss some consequences of the presence of a gap in \\S 3.2.  A similar situation probably takes place in the millisecond pulsar SAX J1808.4--3658, a unique object that demonstrates coherent 2.5~ms flux oscillations (Wijnands \\& Van der Klis 1998; Chakrabarty \\& Morgan 1998; Gilfanov et al. 1998). In this system, matter supposedly flows towards two nearly antipodal magnetic polar caps on the neutron star, and the accretion disk terminates at some distance from the stellar surface. Then, the hot spot located below the disk could be observable. However, the measured nearly sinusoidal pulse profile (Wijnands \\& Van der Klis 1998) suggests that this spot contributes very little, if any, to the total signal. This fact can be used for constraining the geometry and inclination of the system. The upper hot spot should be observable both directly and in X-rays scattered from the disk.  Our scenario may apply to type-I X-ray bursts. Indeed, RXTE detections of nearly coherent oscillations during bursts from eight sources (see Strohmayer 2000 for a review; Wijnands et al. 2000; Galloway et al. 2000) imply the presence of a localized emitting zone. Apparently, this zone slowly evolves in time (on sub-second timescales $\\gg$ the spin period), as the thermonuclear burning front propagates over the stellar surface. It is not clear whether the innermost regions of a standard accretion disk can survive during a burst, when the disk is subject to strong radiation forces. In particular, radiation drag causes the matter flowing from the disk to lose angular momentum and spiral inward (Miller \\& Lamb 1993). Obviously, the disk has a greater chance to remain intact during the rise and decay stages of a burst, when the emitted luminosity, $L$, is much less than the critical Eddington luminosity, $\\ledd$. Note that if the inner disk is destroyed during the maximum of the burst, it can reconstruct very quickly at some later moment when the burst emission has faded substantially, because the characteristic times for all processes taking place within several stellar radii from the star are much shorter than the typical burst duration ($\\sim 10$~s).  Spots of enhanced brightness may also arise in a spreading layer on the surface of a neutron star (Inogamov \\& Sunyaev 1999). This layer rotates differentially, with the azimuthal velocity ranging between the Keplerian velocity ($\\sim 0.5c$) at the equator of the star, and the stellar spin velocity ($\\sim 0.1 c$) near the edges of the layer. Under such circumstances, long-living whirls, analogous to typhoons, Jupiter's famous red spot, or ``cat eyes'', may regularly appear in two bright latitudinal belts within the layer that are located in the upper and lower hemispheres equidistantly from the equator. The velocity shear is largest in these belts. Their location is determined by the accretion rate, namely, they become more distant from the equator (thus closer to the poles) with increasing accretion rate. Also, instabilitites developing in the accretion disk can modulate the accretion rate through the neck of the disk into the spreading layer. In particular, magnetic-field loops originating in the disk due to Balbus-Hawley instability may penetrate into the layer. All these phenomena can lead to the formation of bright emitting zones, rotating faster than the star itself, with quasi-frequencies up to the Keplerian one of $\\sim 2$~kHz (Inogamov \\& Sunyaev 2000; Sunyaev \\& Revnivtsev 2000). If a spot emits a significant fraction of the source luminosity during several spin cycles, quasi-periodic flux oscillations will appear. Averaged (with a fixed period) waveforms of such incoherent oscillations will tend to be nearly sinusoidal. However, the disk-scattered flux should still manifest itself clearly in the power density spectrum of the total signal (i.e., in the ratio of the powers contained in different quasi-harmonics), as well as in the phase delays ($\\lesssim 0.25$) between waveforms measured in different energy bands and in the phase shift ($\\sim 0.25$) of the polarized flux component with respect to total flux.  It is important to note that the considerations above are only applicable to atoll sources, i.e., to LMXBs with luminosities between 0.01 and 0.5 $\\ledd$. In this case, the approximation of a thin accretion disk (Shakura \\& Sunyaev 1973) is appropriate. At luminosities approaching $\\ledd$, like those of Z-sources, the boundary layer extends radially more than one stellar radius and is geometrically thick (Popham \\& Sunyaev 2000). The adjacent disk also becomes thick. Therefore, our picture is no longer valid. Similarly in the case of $L<0.01\\ledd$, Compton cooling becomes inefficient for matter in the disk, and a geometrically thick advection-dominated accretion flow may form (e.g., Narayan et al. 1998). ",
        "conclusions": "We have considered the effect of X-ray scattering by a thin accretion disk on the phase dependence of spin-modulated emission from a neutron star with a weak magnetic field. This effect makes it possible to test the presence of standard thin accretion disks around the neutron stars in LMXBs. The only other existing experimental method to probe Keplerian accretion disks near compact objects (particularly black holes) is through studying the profiles and variability of iron fluorescence lines (see Fabian et al. 2000; Reynolds 2000 for recent reviews). Furthermore, the new method provides a unique opportunity to check whether the neutron star is corotating with the disk, and potentially even to determine the misalignment angle between the spin axis and the disk symmetry axis. Such measurements can significantly enrich our knowledge of the origin and evolution of LMXBs.  The results of this paper demonstrate that the disk-scattered flux component manifests itself better in energy-resolved oscillation waveforms. Nearly coherent oscillations during type-I bursts, especially those detected during the burst rise, are probably the best candidates to look for signatures of disk-scattering. It is possible that RXTE data can already be used for such a search. For example, the phase lag of low-energy photons during a burst from Aql X-1 reported by Ford (1999) may be the result of Doppler boosting in a thin disk. It is important that this interpretation is self-consistent, as opposed to the alternative explanation based on Doppler boosting by the neutron star, as shown in \\S 3.3. The case of the millisecond pulsar SAX~J1808.4--3658, for which similar phase lags and also distortions of the pulse profile have been detected (Cui et al. 1998; Revnivtsev 1999; Ford 2000), appears to be more complicated. The principal difficulty here is that the observed energy spectrum has an approximately power-law shape. If this spectral distribution originates at the hot spot, then no energy dependence of the pulse profile can arise as a result of Doppler boosting (see, e.g., Chen \\& Shaham 1989).  More detailed studies of the waveforms of X-ray flux oscillation in LMXBs could be possible with $\\sim 6$~m$^2$ X-ray detectors (i.e., 10 times larger than RXTE), such as the recently proposed EXTRA and RAE missions (Barret 2000; Kaaret 2000; http://www.cesr.fr/$\\sim$barret/extra.html). A particularly promising way to look for spin-phase signatures of disk-scattering, at least in the millisecond pulsar SAX~J1808.4--3658, is by means of sensitive X-ray polarimetry. The case of SAX~J1808.4--3658 (or similar objects which may be discovered in the future) is favourable because this source demonstrates regular pulsations during long time intervals, which allows one to accumulate a statistically large amount of data. On the contrary, such statistics is hardly achievable for nearly coherent burst oscillations, or for possible quasi-periodic oscillations originating in the spreading layer. Previously, only the time-integrated polarization due to Thomson scattering in the accretion disk (without allowance for its rotation) was discussed as an observational task for near-future X-ray polarimeters (Meszaros et al. 1988). Such polarimeter will be part of the scientific payload of the Spectrum-X-Gamma observatory (Kaaret et al. 1993; http://hea-www.harvard.edu/$\\sim$kaaret/sxrp/index.html). Also, NASA is currently attaching great importance to a programme of long duration balloon flights, lasting up to 100 days (Jones 2000). This programme will no doubt rekindle interest to balloon flights of X-ray polarimeters."
    },
    "0011/astro-ph0011434_arXiv.txt": {
        "abstract": "The Ultra Luminous X-ray Sources (ULXs) are unique in exhibiting moderately bright X-ray luminosities, $L_{\\rm x} \\sim 10^{38-40} {\\rm erg~s^{-1}}$, and relatively high blackbody temperatures, $\\Tin \\sim 1.0-2.0 {\\rm keV}$. From the constraint that $L_{\\rm x}$ cannot exceed the Eddington luminosity, $L_{\\rm E}$, we require relatively high black-hole masses, $M\\sim 10-100 M_\\odot$, however, for such large masses the standard disk theory predicts lower blackbody temperatures, $\\Tin < 1.0$ keV. To understand a cause of this puzzling fact, we carefully calculate the accretion flow structure shining at $\\sim L_{\\rm E}$, fully taking into account the advective energy transport in the optically thick regime and the transonic nature of the flow. Our calculations show that at high accretion rate ($\\dot M \\ga 30~L_{\\rm E}/c^2$) an apparently compact region with a size of $\\Rin \\simeq (1-3)\\rg$ (with $\\rg$ being Schwarzschild radius) is shining with a blackbody temperature of $\\Tin \\simeq 1.8 (M/10M_\\odot)^{-1/4}$ keV even for the case of a non-rotating black hole. Further, $\\Rin$ decreases as $\\dot M$ increases, on the contrary to the canonical belief that the inner edge of the disk is fixed at the radius of the marginally stable last circular orbit. Accordingly, the loci of a constant black-hole mass on the {\\lq\\lq}H-R diagram\" (representing the relation between $L_{\\rm x}$ and $\\Tin$ both on the logarithmic scales) are not straight but bent towards the lower $M$ direction in the frame of the standard-disk relation.  We also plot the ASCA data of some ULXs on the same H-R diagram, finding that they all fall on the regions with relatively high masses, $M \\sim 10-30M_{\\odot}$, and high accretion rates, ${\\dot M}\\ga 10~L_{\\rm E}/c^2$. Interestingly, IC342 source 1, in particular, was observed to move along the constant-$M$ line (not constant $\\Rin$ line) in our simulations. This provides a firm evidence that at least some ULXs are shining at $\\ga L_{\\rm E}$, and containing black holes with $M \\simeq 10-100 M_\\odot$. ",
        "introduction": "Already in the late 1980's, $Einstein$ satellite had found luminous X-ray sources with X-ray luminosities $L_{\\rm x} \\sim 10^{38-40}~{\\rm erg~s^{-1}}$ in nearby spiral galaxies (Fabbiano.\\ 1989). They locate in the off-nucleus region of the host galaxy, hence cannot be an active-galactic nucleus (AGN) of low-luminosity. Only a small number of them are optically identified as young supernova remnants (e.g., the most luminous source in NGC~6946; Schlegel\\ 1994), whereas we lack secure optical identifications for most of them. At present day, these luminous, non-AGN and non-supernova-remnant sources are called {\\lq\\lq}Ultra Luminous X-Ray Sources (ULXs)\" (see Makishima et al.\\ 2000 and reference therein). Apparently, their X-ray luminosities greatly exceed the Eddington limit for a neutron star, $\\sim 2\\times 10^{38}$ erg~s$^{-1}$, by a couple of orders of magnitude. ULXs are thus more like black-hole binaries (BHBs), although we need spectral information to reveal their emission mechanism.  Recently, much progress has been made on the study of ULXs. Thanks to the fine spectral and imaging capability of $ASCA$, high-quality spectra have been accumulated for numbers of ULXs, including NGC~1313 source~A and B (Petre et al.\\ 1994), M81~X-6 (Uno 1996), two ULXs in IC~342 (Okada et al.\\ 1998), and so on. Especially, rather extensive study has been performed by Makishima et al.\\ (2000) and Mizuno et al. (2000).  According to their study, the spectra of most ULXs in nearby galaxies can be well fitted with multi-color disk (MCD) model (Mitsuda et al.\\ 1984), reinforcing its BHB interpretation. However, obtained disk temperatures, typically ranging 1.0--2.0~keV, somewhat contradict the BHB hypothesis. Theoretically, they are too high to be achieved by a standard-type disk around such a high-mass, Schwarzschild (non-rotating) black hole. They are also higher than those observed from well-studied Galactic/Magellanic BHBs (e.g., Tanaka \\& Lewin\\ 1995). The observed high disk temperatures are hence the most severe problem regarding ULXs. To cope with this issue, some authors discussed the possibility of Kerr BH, although their arguments still remain rather qualitative. We need both the detailed calculation of spectra from BHBs and the careful comparison with the observations.  In this $Letter$ we consider an important process which has not been fully appreciated in the study of ULXs. The accretion-disk theory predicts that for high luminosity, comparable to $L_{\\rm E}$, advective energy transport dominates over energy loss by radiation. Abramowicz et al. (1988) were first to construct such a disk model, currently known as the slim disk model (see also Szuszkiewicz et al. 1996; Wang et al. 1999). More recently, Watarai et al. (2000) and Mineshige et al. (2000) discussed the distinctive spectral features of the slim disks in connection with BHBs and narrow-line Seyfert 1 galaxies (NLS1s), respectively (see also Fukue 2000). Following the same line, we calculate the structure and emission properties of high-luminosity accretion disks and demonstrate that the X-ray spectral properties of ULXs are indeed well understood by the slim disk model. In section 2 we give basic equations with advection, where we stress the necessity of solving flow structure with careful consideration of its transonic nature. We then show in section 3 the results of model fitting and compare the theory with ULXs observed by $ASCA$. In section 4, we argue that ULXs do really undergo super-critical accretion.  The final section is devoted to conclusions. Throughout the present study, we normalize $\\dot M$ by $L_{\\rm E}/c^2$ (with $c$ being speed of light); i.e., the normalized mass-flow rate is $ {\\dot m} \\equiv \\dot M/(L_{\\rm E}/c^2) \\simeq \\dot M/ (1.4\\times 10^{17}(M/M_\\odot)~{\\rm g~s}^{-1})$. ",
        "conclusions": "\\noindent{1.} All the ULXs so far observed fall on the regions with high mass ($M \\sim 30M_\\odot$) and high accretion rate ($\\sim 30 L_{\\rm E}/c^2$) on the X-ray H-R diagram (figure 1).  \\noindent{2.} The slim disk model predicts apparent decrease in $\\Rin$ as $L$ increases.  This results in a qualitatively different evolutionary path of a single object on the X-ray H-R diagram, compared with that of the standard disk (i.e., constant $\\Rin$ line). ULXs that are observed at several different epochs (only IC342 was observed at same epochs) actually follow the constant-$M$ lines of the slim disk, not those of the standard disk. This provides a firm evidence that at least some ULXs are near critical accretion phase.  \\noindent{3.} It is not always clear whether a central black hole is rotating or not. What we can conclude is that high luminosity, $\\sim L_{\\rm E}$, is realized in ULXs, not depending on the nature of black holes.  \\noindent{4.} In future work, we need more accurate evaluations of GR effects and the effects of changing viscosity prescriptions. Also, more extensive observational studies, particularly independent estimates of inclination angles, are indispensable for fully understanding the nature of ULXs."
    },
    "0011/astro-ph0011328_arXiv.txt": {
        "abstract": "This contribution describes the difficult task of inferring the IMF from local star-count data, by discussing the mass--luminosity relation, unresolved binary, triple and quadruple systems, abundance and age spreads and Galactic structure, all of which must be accounted for properly for the results to be meaningful. A consensus emerges that the local IMF may be represented by a two-part power-law, with indices $\\alpha=1-1.5$ for stars with mass $m\\simless 0.5\\,M_\\odot$, and the Salpeter value $\\alpha=2.3$ for more massive stars, but some uncertainties remain. Notable is also that the sensitivity of the stellar luminosity function (LF) to the derivative of the mass--luminosity relation is very evident in the (local) Hipparcos and HST, open-cluster and globular-cluster LFs, thus allowing tests of stellar structure theory.  The upcoming astrometry space missions DIVA and GAIA will undoubtedly lead to significant advances in this field. ",
        "introduction": "\\noindent The distribution of masses of stars born together, the {\\it initial mass function} (IMF), determines the appearance and evolution of galaxies and star clusters, and composes the boundary condition for star-formation when it ends.  Any affirmed variation of the IMF with local conditions would pose tremendously important constraints on our understanding of how stars form.  The detailed shape of the IMF, let alone any variation of it between populations, has been arduous to distill from the observational data, with a large amount of serious work remaining to be done in this fundamentally important field.  This text focuses on the derivation of the IMF from solar-neighbourhood (sn) star-count data.  The sn constitutes a particularly important stellar sample because of it's proximity, and because it represents a mixture of many star-formation (sf) events. The {\\it solar-neighbourhood IMF} is therefore an {\\it average IMF}, being valid for the on average about~5~Gyr old Milky-Way (MW) disk stars.  Furthermore, by studying the sn, methods can be developed that are of general use for the interpretation of more distant stellar populations.  Because of the restriction to the sn, the mass range covered by this contribution spans $0.1-3\\,M_\\odot$, more massive stars being too rare. Very-low-mass stars and brown dwarfs are dealt with by Chabrier \\& Baraffe (2000), while Massey (1998) addresses the IMF for massive stars. Empirical evidence for systematic variations over all stellar masses ($10^{-2}-10^2\\,M_\\odot$) is studied by Kroupa (2000a), and Meyer et al. (1999) review the IMF in star-forming regions. ",
        "conclusions": ""
    },
    "0011/astro-ph0011091_arXiv.txt": {
        "abstract": "A small fraction of double radio sources show a peculiar and striking hybrid morphology; they have a distinctly FR I  structure on one side of the nucleus, and a FR II structure on the other.  We argue that the mere existence of these HYMORS is quite incompatible with the theoretical explanations for the Fanaroff--Riley dichotomy that are based upon the nature of the jet plasma, or those invoking an intrinsic property of the central engine.  Rather, these HYMORS strongly support models that explain the difference between FR I and FR II sources in terms of asymmetry of interaction of the jets with the external environments.  We further show that a model for radio source dynamics we had earlier proposed can neatly reproduce the observed dependence of the radio power dividing the two FR classes on the optical luminosity of the host galaxy, as found by Owen \\& White and Ledlow \\& Owen. ",
        "introduction": "The vast majority of extragalactic radio sources can be easily classified into the more powerful, edge-brightened FR II sources and the weaker, edge-dimmed FR I sources (Fanaroff \\& Riley 1974).  The boundary between these classes in terms of monochromatic radio power was shown to rise with increasing optical luminosity of the host galaxy: $P^*_R \\propto L_{\\rm opt}^{1.65}$ (e.g., Ledlow \\& Owen 1996, Fig.\\ 1).  The large number of explanations for the Fanaroff-Riley dichotomy can be divided into {\\it intrinsic} models, where some property of the central engine launches entirely different types of jets, and {\\it extrinsic} models, where very similar jets are ejected in both FR classes but interactions with the environment determine their large-scale morphologies (Gopal-Krishna \\& Wiita 2000a --- GKW00, and references therein).  Among the better developed intrinsic models are those where: faster spinning black holes launch FR II jets and FR I jets come from slower spinning BHs; jet plasma is composed of $e^-$--$p$ for FR II's  and $e^-$--$e^+$ for FR I's; advection dominated accretion disks yield FR I's while standard thin disks yield FR II's. Possible support for such models comes from the differences between the properties of the host galaxies of the two FR classes (e.g., Baum et al.\\ 1995).  Various extrinsic models all argue that the jets differ, if at all, only in their power or thrust, so the morphological dichotomy arises from the differences in the interactions of the jets with the gaseous media. Weaker and/or poorly collimated jets, or those ploughing through a more disruptive environment, quickly lose their terminal hot-spots (Gopal-Krishna \\& Wiita 1988), or are subject to instabilities that dramatically increase the entrainment of ambient gas (e.g., De Young 1993; Bicknell 1984, 1995), both processes leading to a FR I morphology. Probably the strongest evidence for the basic similarity of the jets comes from the VLBI observations showing that the jets in most FR I, as well as FR II, radio sources start off with relativistic bulk velocities; however, only FR II jets are able to remain relativistic to very large distances (e.g., Laing et al.\\ 1999 and references therein). ",
        "conclusions": ""
    },
    "0011/astro-ph0011058_arXiv.txt": {
        "abstract": "One of the key questions to understanding the efficiency of diffusive shock acceleration of the cosmic rays (CRs) is the injection process from thermal particles. A self--consistent injection model based on the interactions of the suprathermal particles with self--generated magneto--hydrodynamic waves has been developed recently by Malkov (\\cite{Malk98}). By adopting this analytic solution, a numerical treatment of the plasma--physical injection model at a strong quasi--parallel shock has been devised and incorporated into the combined gas dynamics and the CR diffusion--convection code. In order to investigate self--consistently the injection and acceleration efficiencies, we have applied this code to the CR modified shocks of both high and low Mach numbers ($M=30$ and $M=2.24$) with a Bohm type diffusion model. Both simulations have been carried out until the maximum momentum $(p_{\\rm max}/m_{\\rm p}c) \\sim 1$ is achieved to illustrate early evolution of a Bohm type diffusion. We find the injection process is self--regulated in such a way that the injection rate reaches and stays at a nearly stable value after quick initial adjustment. For both shocks about $10^{-3}$ of the incoming thermal particles are injected into the CRs. For the weak shock, the shock has reached a steady state within our integration time and $\\sim 10\\%$ of the total available shock energy is transfered into the CR energy density. The strong shock has achieved a higher acceleration efficiency of $\\sim 20 \\%$ by the end of our simulation, but has not yet reached a steady--state. With such efficiencies shocks do not become CR--dominated or smoothed completely during the early stages when the particles are only mildly relativistic. Later, as the CR pressure becomes dominated by highly relativistic particles that situation should change, but is difficult to compute, since the maximum CR momentum increases approximately linearly with time for this model. In the near future we intend to extend such shock simulations as these to include much higher CR momenta using an adaptive mesh refinement technique currently under development. ",
        "introduction": "\\label{intro} The non--thermal energy distributions of cosmic ray ions or source distributions of electrons emitting synchrotron radiation in various astrophysical objects are commonly described as produced by the first order Fermi acceleration process  at shocks (for reviews see Drury~\\cite{Drur83}; Blandford \\& Eichler~\\cite{BlEi87}; Kirk et al.~\\cite{KiMePr94}).  When particles diffuse\\footnote{In general, also anomalous transport like {\\em sub}--diffusion or {\\em super}--diffusion can be realized.} off the moving scattering centers in a region divided by a velocity discontinuity (shock), these particles can be accelerated if their mean free paths exceed the shock thickness. The relative momentum gain for a cycle of two crossings of the shock is then proportional to the velocity difference across the shock, i.e.~of first order with respect to the shock velocity (Bell~\\cite{Bell78}). In astrophysical collisionless plasmas an electro--magnetic field must be present to change the energy of particles. Waves or irregularities in this field provide particle scattering, which leads to diffusion. Consider a shock with velocity $u_{\\rm s}>0$ propagating into a plasma at rest with density $\\rho$ and with a homogeneous magnetic field $B_0$ in the direction of the shock normal. The plasma is compressed to the density $\\rho_{\\rm d}$ by the shock, and flows downstream with the velocity $u_{\\rm d}=u_{\\rm s}(1-1/r)$, where $r= \\rho_{\\rm d}/\\rho$ is the compression ratio. Particles with the mean downstream velocity $\\langle v \\rangle =u_{\\rm d}$ cannot cross the shock from downstream to upstream, because $\\langle v \\rangle < u_{\\rm s}$. In addition the shock may not be a discontinuity for a particle at this energy, because the gyro radius of the thermal particles is of the order of the shock thickness, leading to adiabatic energy change while crossing the shock. Because the plasma is also heated downstream of the shock, some supra--thermal particles in the high energy tail of the Maxwellian velocity distribution may gain energies and have velocities that allow them to re--cross the shock. In a homogeneous magnetic field, due to lack of scattering centers these particles would escape upstream without returning to the shock. Then, the acceleration mechanism would not apply. However, the population of particles that can move upstream provide a seed particle beam, which generates Alfv\\'en waves responsible for scattering and, therefore, diffusion, which is an essential element of the first order Fermi acceleration process.  The problem of particle acceleration from thermal energies up to relativistic particle energies is highly non--linear, as first pointed out by Eichler~(\\cite{Eich79}). First, the energy transferred from the bulk of the plasma to the sub--population of accelerated particles can change the thermodynamic properties of the plasma like the temperature and density. In addition, the accelerated particles provide their own pressure in the system, which, since it differs from the thermal pressure, modifies the velocity structure of the shock transition. Second, the waves generated by particles escaping upstream determine the transport properties of the plasma, and, therefore, regulate this wave generating escape itself. The manner in which the wave--particle interactions control the fraction of plasma particles that can escape upstream to participate in the Fermi process is commonly called {\\em injection}. This is a basic aspect of the plasma of collisionless shocks and is itself highly non--linear. This injection problem is fundamentally related to the question of the efficiency of particle acceleration at shocks by the Fermi process.  Different numerical methods have been used to treat the injection problem of CR modified shocks. In Monte--Carlo simulations of non--linear particle acceleration the details of a posited scattering law provide an injection parameter, but one not determined self--consistently from the particle wave interaction (e.g.~Ellison et al.~\\cite{ElBaJo96}; Baring et al.~\\cite{BaElReGrGo99}). In contrast to pure kinematical effects from shock velocity, particle speed and inclination angle of magnetic field and shock, the waves responsible for particle scattering depend on plasma properties like temperature and the beam strength of the wave generating particles itself. These coupled and time dependent effects are not easy to incorporate into a Monte--Carlo approach. Time dependent Monte--Carlo simulations have been presented by Knerr et al.~(\\cite{KnJoEl96}), but still with a prescribed scattering law, as a parameterization of injection.  In the two--fluid approach the cosmic rays are treated as a diffusive gas without following their momentum distribution. The energy transfer into CRs in these models is based on a fraction of the upstream gas particles, that are instantaneously accelerated at the shock (Dorfi~\\cite{Dorf90}), around the shock (Jones \\& Kang~\\cite{JoKa90}) or at velocity gradients (Zank et al.~\\cite{ZaWeDo93}). In practical terms, because the shock is the most prominent velocity gradient in the system, these techniques are very closely related, as pointed out by Kang \\& Jones~(\\cite{KaJo95}). Essentially the same parameterization is also used in the numerical solution of the hydro--dynamical equations coupled to the momentum dependent cosmic--ray transport equation (e.g.~Falle \\& Giddings~\\cite{FaGi87}; Kang \\& Jones~\\cite{KaJo91}; Berezhko et al.~\\cite{BeYeKs94}).  Kang \\& Jones~(\\cite{KaJo95}) used a numerical injection model with two essentially free parameters which describe boundaries in momentum at which particles can be accelerated ($p_1 = c_1 \\cdot \\pth$ where $\\pth$ is the peak momentum of the Maxwell distribution) and from which these contribute to the cosmic--ray pressure ($p_2 = c_2 \\cdot \\pth$). Still, these momentum boundaries are free parameters, which can be translated into a particle fraction of the upstream gas. However, models that incorporate more details of the plasma physics of the background plasma, and, therefore, the CRs at injection energies are really necessary to constrain the phase--space function and therefore determine these parameters. In this way we can incorporate self--consistent plasma physical models into numerical simulations of CR acceleration.  Such a plasma physical model based on non--linear interactions of particles with self--generated waves in a shocked plasma has been investigated numerically by solving the kinetic equations of ions in a magnetic field, and treating the electrons as a background fluid (Quest~\\cite{Ques88}). These simulations show that ions can be scattered back and forth across the shock by self--generated waves, and Quest~(\\cite{Ques88}) also points out that these scattered ions can provide a seed population of cosmic rays.  Recently, the kinetic equations of ions were solved analytically for non--linear wave--fields near strong parallel shocks by Malkov \\& V\\\"olk (\\cite{MaVo95}) and Malkov (\\cite{Malk98}). These authors were able to constrain the fraction of phase--space of the background plasma that can be injected into the acceleration process as a result of the self--regulating interaction between wave generation and particle streaming. Here we incorporate this self--consistent analytical result in numerical solutions of the hydro--dynamical equations together with the cosmic--ray transport equation. Our simulations, therefore, provide the first time dependent solution of the problem of modified shocks that includes a self--consistent plasma physical injection model. This technique enables us to determine the level of shock modification and acceleration efficiency in an evolving shock without a free parameter for the injection process (Gieseler et al.~\\cite{GiJoKa99}).  We describe the plasma physical injection model in some more detail in Sect.~\\ref{model} before we outline the coupled set of dynamical equations of the plasma and cosmic rays in Sect.~\\ref{dyn_model}, together with details of the numerical method we used to solve these equations. The results are presented in Sect.~\\ref{results} and Sect.~\\ref{resultsweak}, consisting of the time dependent evolution of the plasma properties, showing especially the modification of the velocity profile and the momentum distributions of thermal and relativistic cosmic--ray particles. In addition we present our results in the form of a particle injection efficiency and an energy transfer efficiency at modified strong shocks. ",
        "conclusions": "We have developed a numerical method to include self--consistently the injection of the supra--thermal particles into the cosmic--ray population at quasi--parallel shocks according to the analytic solution of Malkov (\\cite{Malk98}). Toward this end, we have adopted the ``transparency function'' $\\tauesc(v,u_2)$ which expresses the probability that supra--thermal particles at a given velocity can leak upstream through the magnetic waves, based on non--linear particle interactions with self--generated waves. We have incorporated the transparency function into the existing numerical code which solves the cosmic--ray transport equation along with the gas dynamics equations. In order to investigate the interaction of high energy particles, accelerated by the Fermi process, with the underlying plasma flow {\\em without} using a free parameter for the injection efficiency, we have applied our code with the new injection scheme to both strong ($M=30$) and weak ($M=2.24$) parallel shocks.  The main conclusions from the simulation results are as follows:  \\begin{enumerate}  \\item{} The injection process is regulated by the overlap of the population of supra--thermal particles in the injection pool and the function ${\\partial \\tauesc(p,t)}/{\\partial p}$. As being in the high energy tail of the Maxwell velocity distribution, the population in the injection pool depends strongly on the gas temperature and the particle momentum. The function ${\\partial \\tauesc(p,t)}/{\\partial p}$ behaves like a delta--function defined near a narrow injection pool. As the postshock gas cools due to high initial injection, the Maxwell distribution shifts to lower momenta. But the transparency function also shifts to lower momenta, as well, due to its dependence on the postshock flow velocity. As a result, the injection rate reaches and stays at a stable value after a quick initial adjustment, and also depends only weakly on the initial conditions. This self--regulated injection may imply a broad application of our simulation methods.  \\item{} The fraction of the background particles that are accelerated to form the non--thermal part of the distribution turns out to be in the range $1.2\\cdot 10^{-3} \\la \\xi^* \\la 1.9\\cdot 10^{-3}$ for the range of initial wave--amplitudes $0.3\\le\\ep\\le 0.4$ at a $M=30$ shock. For a $M=2.24$ shock, a slightly higher injection is achieved at $\\xi^*= (2.5\\pm 0.7)\\cdot 10^{-3}$, but this could be a lower limit. Such values for the particle injection efficiency have been used as a parameter for spherically expanding SNRs by several authors (Dorfi~\\cite{Dorf90}; Jones \\& Kang~\\cite{JoKa92}; Berezhko et al.~\\cite{BeKsYe95}; Berezhko \\& V\\\"olk~\\cite{BeVo00}). These values are well above the ``critical injection rate'' of $\\eta_{\\rm crit} \\sim 10^{-4}$ above which spherical shocks of this Mach number are CR dominated according to Berezhko et al.~(\\cite{BeKsYe95}).  \\item{} Due to computational limitations of using a Bohm type diffusion model, we have integrated our models until the maximum momentum reaches about $(p_{\\rm max}/m_{\\rm p}c) \\sim 1$. For the $M=30$ shock model, the energy flux in the total CR distribution was about $18\\%\\pm5\\%$ of the energy flux in the thermal plasma and shocks didn't become CR dominated and smoothed completely by the end of our simulations. For the $M=2.24$ shock model, the acceleration efficiency is lower by a factor of two compared to the high Mach shock because of the smaller velocity jump across the shock.  \\item{} Just above the injection pool, the distribution function changes sharply from a Maxwell distribution to an approximate power--law whose index is close to the test--particle slope. We estimated this critical momentum as $\\pinj \\sim (2.2-2.3) \\cdot \\pth$ where $\\pth = 2\\sqrt{ m_{\\rm p} \\boltz T}$. This determines the number of particles in the injection pool by $f(\\pinj) \\propto \\exp ( - \\pinj^2 / 2 m_{\\rm p} \\boltz T ) $. For strong shocks this translates into a distribution function at injection energies of $g(\\pinj) \\sim (1/100 - 1/200) g(\\pth)$.  \\end{enumerate}  While the weak shock model of $M=2.24$ reaches a steady--state, the strong shock model of $M=30$ has not reached a steady--state by the end of our simulation. We expect for the strong shock that the CR pressure continues to increase and the shock becomes CR dominated, leading to the greater total velocity jump and more efficient acceleration. In realistic shocks such as SNRs, however, escaping particles due to non--planar geometry or lack of scattering at high momentum are likely to become important. To resolve this non--linear evolution, much longer physical time scales have to be simulated, until CRs reach energies where escape is likely to be important. The key problem here is the range in configuration and momentum space that has to be computed. Our method uses a grid with uniform cells in configuration space, chosen fine enough to capture the evolution of $g(x,p)$ at near--thermal momenta where the diffusion coefficient is proportional to $p^2$ (Bohm diffusion). This leads to a computationally extremely expensive calculation, especially because the grid has to be large enough to contain the diffusion length scale of the highest momentum CRs. The problem can be solved on a much larger time scale by using an adaptive mesh refinement (AMR) code with the shock tracking techniques (Kang \\& Jones~\\cite{KaJo99}). In the near future we plan to incorporate the injection model presented here into the powerful shock tracking AMR--code, to calculate the evolution of the phase--space distribution of the plasma during different phases of SNRs. This would allow us to investigate with a plasma--physical based injection model how the slowly growing cosmic--ray pressure at a strong shock eventually modifies the shock structure. A strong modification will cause the velocity jump across the subshock to decrease and the distribution function of the suprathermal particles to steepens. This might have further back reaction on the injection efficiency. Also the CR distribution will deviate from a simple power--law. For a calculation up to the highest energy CRs, also the spherical geometry of a SNR should be taken into account. Such an approach could lead to a consistent calculation of the complete phase--space distribution at quasi--parallel shocks, and should be a promising step towards a calculation of the overall efficiency of SNRs in producing CRs during their evolution.  For oblique shocks, the injection efficiencies calculated here for a parallel shock should define an upper limit, because the statistical probability of a particle to cross the shock from downstream to upstream decreases with the intersection velocity of magnetic field and shock front. This kinematical effect was investigated by Baring et al.~(\\cite{BaElJo93}) with the use of Monte--Carlo simulations. However, in the model we have incorporated here, the injection is already suppressed strongly (compared to the purely kinematical model) by the reduced transparency of the plasma due to the high amplitude Alfv\\'en waves. We point out, that for oblique shocks, the filtering due to Alfv\\'en waves may be reduced due to the decreased downstream amplification of the wave amplitude. This would allow lower energy particles to be injected, and the kinematical effect could be partly balanced. As a result, we speculate that the dependence on the obliquity might be significantly weaker than calculated by Baring et al.~(\\cite{BaElJo93}). Resolution of that important question must await more complete  understanding of the injection physics.  In summary, we have shown that the process of particle acceleration under consideration of a plasma physical injection model underlies a rather effective self--regulation. Apart from the direct particle--wave interaction described by the injection model itself, also the energetic feedback of the energy transfer between thermal plasma and cosmic--rays keeps the fraction of particles in the non--thermal distribution at roughly $10^{-3}$ of the particles swept through the shock. These self--regulation mechanisms lead to a quite stable injection efficiency, which depends weakly on the initial conditions. \\vspace*{-0.3cm}"
    },
    "0011/astro-ph0011235_arXiv.txt": {
        "abstract": " ",
        "introduction": "The probability that high-redshift quasars are gravitationally-lensed by intervening galaxies increases rapidly with the cosmological constant, $\\Omega_{\\Lambda_0}$ (whilst being only weakly dependent on the density parameter, $\\Omega_{{\\rm m}_0}$), and the low number of lenses observed implies that $\\Omega_{\\Lambda_0} \\la 0.7$ (e.g.\\ Kochanek 1996). One of many uncertainties has been the (small) core radii of elliptical galaxies, which, at least naively, reduce their lensing cross-section. However, if ellipticals are normalised relative to their observed line-of-sight velocity dispersion, $\\sigma_{||}$, then increasing the core radius must result in an increased mass normalisation (specified by the assymptotic velocity dispersion, $\\sigma_\\infty$). ",
        "conclusions": "The use of self-consistent dynamical models for elliptical galaxies shows that, independent of their core radius, they are effective lenses. The results imply that the current upper limits placed on $\\Omega_{\\Lambda_0}$ due to the low observed frequency of quasar lenses are not weakened by any uncertainty in the core structure of ellipticals -- in fact the existing limits are probably more robust than previously believed."
    },
    "0011/gr-qc0011044_arXiv.txt": {
        "abstract": "In association with the Blanford-Znajek mechanism for rotational energy extraction from Kerr black holes, it is of some interest to explore how much of magnetic flux can actually penetrate the horizon at least in idealized situations. For completely uncharged Kerr hole case, it has been known for some time that the magnetic flux gets entirely expelled when the hole is maximally-rotating. In the mean time, it is known that when the rotating hole is immersed in an originally uniform magnetic field surrounded by an ionized interstellar medium (plasma), which is a more realistic situation, the hole accretes certain amount of electric charge. In the present work, it is demonstrated that as a result of this accretion charge small enough not to disturb the geometry, the magnetic flux through this slightly charged Kerr hole depends not only on the hole's angular momentum but on the hole's charge as well such that it never vanishes for any value of the hole's angular momentum. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011003_arXiv.txt": {
        "abstract": "Aspherical explosion models for the hypernova (hyper-energetic supernova) SN~1998bw are presented. Nucleosynthesis in aspherical explosions is examined with a two-dimensional hydrodynamical code and a detailed nuclear reaction network. Aspherical explosions lead to a strong $\\alpha$-rich freezeout, thus enhancing the abundance ratios [$^{44}$Ca, $^{48}$Ti, and $^{64}$Zn / Fe] in the ejecta. The nebular line profiles of the Fe-dominated blend near 5200 \\AA\\ and of [\\OI] 6300,6363 \\AA\\ are calculated and compared with the observed late time spectra of SN 1998bw. Compared with the spherical model, the unusual features of the observed nebular spectra can be better explained if SN 1998bw is a strongly aspherical explosion with a kinetic energy of $\\sim 10^{52}$ ergs viewed from near the jet direction. ",
        "introduction": "\\label{sec:int}  The exceptionally bright Type Ic supernova (SN Ic) SN 1998bw was discovered as the probable optical counterpart of the gamma-ray burst GRB980425 (Galama et al. 1998). The early light curve and the spectra of SN 1998bw have been successfully modeled as the hyper-energetic explosion (kinetic energy $E \\sim$ 4 $\\times$ 10$^{52}$ ergs) of a massive C+O star (Iwamoto \\etal 1998; Woosley \\etal 1999; Branch 2001). In this paper the term \"hypernova\" is used to refer to a SN explosion with $E \\gsim 10^{52}$ ergs, regardless of the nature of the central engine (Nomoto et al. 2001a).  Despite the success of the hypernova model in reproducing the observed features of SN 1998bw at early times, some properties of the observed light curve and spectra at late times are difficult to explain. (1) The tail of the observed light curve declines more slowly than the synthetic curve, indicating that at advanced epochs $\\gamma$-ray trapping is more efficient than expected (Nakamura \\etal 2001a; Sollerman \\etal 2000). (2) In the nebular epoch, the [\\OI]6300\\AA\\ emission is narrower than the emission near 5200\\AA. As discussed in Mazzali \\etal (2001), this latter feature is mostly due to a blend of [\\FeII] lines. Mazzali \\etal (2001) calculated synthetic nebular-phase spectra of SN~1998bw using a spherically symmetric NLTE nebular code based on the deposition of $\\gamma$-rays from $^{56}$Co decay in a nebula of uniform density and composition. They showed that the [\\OI] and the [\\FeII] features can only be reproduced if different velocities are assumed for the two elements. A significant amount of slowly-moving O is therefore necessary to explain the profile of the [\\OI] line.  Both these features are in conflict with what is expected from a spherically symmetric explosion  model, where $\\gamma$-ray deposition decreases with time and where iron is produced in the deepest layers and thus has a lower average velocity than oxygen. Mazzali \\etal (2001) suggested that these are signatures of asymmetry in the ejecta.  Therefore in this paper we examine aspherical explosion models for hypernovae.  Aspherical explosions of massive stars have been investigated as possible sources of gamma-ray bursts (GRBs) (Woosley 1993; Paczynski 1998). MacFadyen \\& Woosley (1999) showed numerically that the collapse of a rotating massive core can form a black hole with an accretion disk, while a jet emerges along the rotation axis. The jet produces a highly asymmetric explosion (Khokhlov \\etal 1999). However, these studies did not calculate explosive nucleosynthesis, nor did they show spectroscopic and photometric features of aspherical explosions. Nagataki (2000) performed nucleosynthesis calculations for aspherical SN explosions to explain some features of SN~1987A, but he only addressed the case of a normal explosion energy.  In the present study, we examine the effect of aspherical explosions on nucleosynthesis in hypernovae.  We then investigate the degree of asphericity in the ejecta of SN 1998bw, which is critically important information to confirm the SN/GRB connection, by computing synthetic spectra for the various models viewed with different orientations and comparing the results with the observed late time spectra of SN~1998bw (Patat \\etal 2001). ",
        "conclusions": "\\label{sec:con}  We calculated the nucleosynthesis in aspherical hypernova explosions. We find that in such explosions Fe is mostly ejected at high velocity in a jet along the polar direction, while nearer the equatorial plane burning is less effective, and low-velocity O is mostly ejected. We show that the unusual ratio of the width of the O and Fe nebular lines in SN~1998bw can be explained with a strongly aspherical explosion model viewed from a near-jet direction. Also, in this case the O line has a very sharp peak, in agreement with the observations.  Much of our discussion was based on the identification of the 5200\\AA\\ feature as a blend of [\\FeII] lines. Although several caveats apply to that identification, because other lines may contribute, we claim that [\\FeII] lines dominate. Other possible contributions are as follows. According to Sollerman et al. (2000) the feature contains not only [\\FeII] lines but also lines of [\\MgI], [\\OI] and possibly [\\FeI].  Lines of [\\MgI] are included in our nebular code, and the strength of the 5172\\AA\\ emission, which is consistent with the 5470\\AA\\ line reproducing the observed [\\MgI] peak near 4500\\AA\\ is much smaller than that of the [\\FeII] lines. [\\OI] 5577\\AA\\ is also included in our code, and it is strong and fills up the red part of the emission, but it does not contribute to the blue side (Mazzali et al. 2001).  As for [\\FeI], Sollerman et al. (2000) say that their models `have too low a degree of ionization', suggest that the density in the Fe-emitting regions should be lower, and `regard the FeI emission as dubious'. We confirm this result.  Another possibility is allowed \\FeII\\ emission, such as identified by Filippenko (1989) in the Type II SNe 1987F and 1988I and by Filippenko \\etal (1990) in the Type Ic SNe 1985F and 1987M. The main \\FeII\\ features are at 4570\\AA\\ (multiplets 37 and 38), 5190\\AA\\ (multiplet 42) and 5320\\AA\\ (multiplets 42 and 43). These lines are not included in our atomic model, because not all the collision strengths for these transitions are available. In particular multiplet 42 is not available. However, several arguments apply against the \\FeII\\ identification in SN~1998bw. 1) In SNe 1987F and 1988I all three features are strong, while in the SNe 1985F and 1987M the feature at 4570\\AA\\ is narrow and it is identified as [\\MgI], suggesting that the feature at 5200\\AA\\ could be [\\FeII]. A similar situation holds for SN~1998bw, also a SN~Ic. 2) \\FeII\\ emission occurs at high density. Filippenko (1989) infers $\\log n_e \\sim 9$ -- 10 for SN~1987F. However, the densities we derive for SN~1998bw are more than one order of magnitude smaller (Mazzali \\etal 2001). 3) The relatively low density is confirmed by the large ratio ($\\sim 4$) of the \\CaII\\ IR emission compared to [\\CaII] 7300\\AA. Fig.2 of Ferland \\& Persson (1989) suggests that $\\log n_e < 9$. Finally, the models of Mazzali \\etal (2001), which are based on [\\FeII] only, reproduce the 5200\\AA\\ feature very well.  In any case, extending the model to include \\FeII\\ lines is a worthwhile effort which we are going to make.  The 5200\\AA\\ feature is bluer on Nov. 26 and thereafter than on earlier epochs. Two principal factors are probably responsible for that: 1) the disappearance of the continuum, which is still significant in the earlier spectra; 2) the reduced intensity of \\OI] 5577\\AA, as shown also in our models (Mazzali et al. 2001). This is a high-density line, and it decreases quickly in strength as the ejecta expand. Fading \\FeII \\ emission may also play a role in causing a small wavelength shift.  Weak unaccounted emission is present to the blue of the [\\FeII] feature (4800-5100\\AA). A similar emission is present in the nebular spectra of SNe~Ia as well (Patat \\etal 2001), and it is poorly reproduced by synthetic spectra. It may be due to a forest of weak [\\FeII] and [\\CoII] transitions whose atomic parameters are not well known.  The [\\OI] line in other SNe~Ib/c (Filippenko \\etal 1995, Matheson et al. 2001) also shows a strongly peaked profile, as in SN~1998bw. This probably signals the existence of oxygen at low velocity in most SNe Ib/c. Although there are very few SNe Ib/c in which the feature centered at 5200\\AA\\ is detected with enough S/N, it appears that in these objects the [\\OI] line has a similar profile. For low energy supernovae ($1 \\times 10^{51}$ergs), the spherical model G shows a sharply peaked [\\OI] line profile (Fig. 7). This seems to favor spherical explosions for these low energy supernovae. However, it is premature to conclude that asymmetry is completely absent in low energy SNIb/c, because the effects of asphericity in a late time spectra is not so large in case of normal supernovae as seen in Figure 7. This is due to the low expansion velocities. The oxygen line can be narrow even in the case of an aspherical explosion viewed from near the equator. For example, it is difficult to distinguish model D viewed from 75 deg from model G. Because the probability of viewing an asphericl SN from the near polar direction is smaller than it is to view it near the equator, the observations are not inconsistent with the possibility that most SNe Ib/c are more or less aspherical. More nebular spectra of SNe Ib/c with higher S/N ratio are needed for further investigation.  Our aspherical explosion models may be able to explain the slow decline of the late light curve of SN 1998bw. In these models, the equatorial region is denser than in a spherically symmetric model with the same explosion energy. At advanced epochs this region may be able to trap $\\gamma$-rays more efficiently than a spherical model, as first suggested by Nakamura et al. (2001a). Chugai (2000) showed that a spherical model could reproduce the light curve, if the density near the center (i.e. in the Fe-dominated region) was increased above that of the hydrodynamical model of Iwamoto \\etal (1998). However, Sollerman et al (2000) find that the O-dominated region should be dense, and the Fe-dominated region not dense. Although taken individually these conclusions are probably correct, they appear to be in conflict with one another. An aspherical model offers a natural solution, because it predicts the presence of high density, O-dominated matter near the center (Fig. 3).  A small degree of linear optical polarization ($\\sim 0.5$\\%) was reported in SN 1998bw (Kay \\etal 1998; Iwamoto \\etal 1998; Patat \\etal 2001). This can be explained with different combinations of asphericity and viewing angle, or with large scale clumping in a basically spherical envelope. One solution is that ejecta with a moderate departure from sphericity are viewed from slightly off the axis of symmetry. Our strongly aspherical explosion model has an axis ratio of about 3:2 at the outer edge of the oxygen envelope (Fig. 3), therefore it is consistent with the observed polarization, if it is viewed from near the jet direction.  \\bigskip  This work has been supported in part by the grant-in-Aid for COE Scientific Research (07CE2002, 12640233) of the Ministry of Education, Science, Culture, and Sports in Japan."
    },
    "0011/astro-ph0011529_arXiv.txt": {
        "abstract": "The recent discoveries of luminous quasars at high redshifts imply that black holes more massive than a few billion solar masses have been assembled already when the universe was less than a billion years old.  We show that the existence of these black holes is not surprising in popular hierarchical models of structure formation. For example, the black hole needed to power the quasar SDSS 1044-0125 at $z=5.8$ can arise naturally from the growth of stellar-mass seeds forming at $z>10$, when typical values are assumed for the radiative accretion efficiency ($\\sim0.1$), and the bolometric accretion luminosity in Eddington units ($\\sim1$).  Nevertheless, SDSS 1044-0125 yields a non--trivial constraint on a combination of these parameters. Extrapolating our model to future surveys, we derive the highest plausible redshift for quasars which are not lensed or beamed, as a function of their apparent magnitude.  We find that at a limiting magnitude of $K\\sim 20$, quasar surveys can yield strong constraints on the growth of supermassive black holes out to $z\\sim 10$. ",
        "introduction": "Quasars have long been believed to be powered by the accretion of gas onto supermassive black holes (Salpeter 1964; Zel'dovich 1964; Lynden-Bell 1969).  Dynamical studies indicate that remnant black holes (BHs) indeed reside in the quiescent nuclei of most nearby galaxies (e.g. Magorrian et al. 1998; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000, and references therein), and imply that BH formation is a generic consequence of galaxy formation.  Although the activity of bright quasars peaks at $z\\sim 2.5$, quasars are known to exist at higher redshifts.  Approximately 200 quasars have been found so far at $z>4$ (the majority of which are in the Palomar Sky Survey\\footnote{See http://astro.caltech.edu/\\~{}george/z4.qsos for a comprehensive list of known $z>4$ quasars.}), including the record--holding bright quasar SDSS 1044-0125 at $z=5.80$, recently discovered by the Sloan Digital Sky Survey (SDSS; Fan et al. 2000).  Supermassive BHs can grow out of low-mass ``seed'' BHs through accretion or mergers (Rees 1984; Barkana \\& Loeb 2000a).  In popular hierarchical models for structure formation, the first collapsed gaseous objects have low masses.  It is natural to identify these galaxies with baryonic masses just above the cosmological Jeans mass, $\\sim 10^{4}~{\\rm M_\\odot}$, as the first sites where seed BHs may form (Larson 2000).  Here we postulate that supermassive BHs are the merger products of individual BHs that grow out of these seeds through gas accretion.  The natural $e$--folding timescale for the growth of a single seed can be written as \\begin{equation} t_{\\rm acc}\\equiv {M_{\\rm bh}\\over \\dot M_{\\rm bh}} = 4 \\times 10^7 \\left({\\epsilon \\over 0.1}\\right) \\eta^{-1}~{\\rm yr}, \\label{eq:1} \\end{equation} where $\\epsilon\\equiv L_{\\rm bol}/\\dot M_{\\rm bh}c^2$ is the radiative efficiency for a mass accretion rate $\\dot M_{\\rm bh}$, and $\\eta\\equiv L_{\\rm bol}/L_{\\rm Edd}$ is the bolometric accretion luminosity in Eddington units, $L_E=4\\pi GM_{\\rm bh} c \\mu_e m_p/\\sigma_T$. Here $\\sigma_T$ is the Thomson cross-section and $\\mu_e=1.15$ is the mean atomic weight per electron.  The growth of a $10^9M_\\odot$ BH out of a stellar mass seed requires about $\\ln (10^9M_\\odot/10M_\\odot)=18.4$ $e$--folding times or $\\sim 7\\times 10^8 (\\epsilon/0.1)\\eta^{-1}$ yrs. This time is shorter than the age of the universe at $z=5.8$ if $(\\epsilon/0.1)\\eta^{-1}\\la 1$. However, the time available for the growth of galaxies is shorter than the entire age of the universe at the corresponding redshift.  Turner (1991) recognized that exceptionally bright, high--redshift quasars may yield interesting constraints on cosmic structure formation.  Motivated by the recent discovery of the SDSS 1044-0125 quasar at $z=5.80$ and the emergence of a concordance cosmological model for structure formation\\footnote{Throughout this {\\it Letter} we adopt a $\\Lambda$CDM cosmology with $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_{\\rm b}=0.045$, $h=0.7$, $\\sigma_{8h^{-1}}=0.9$, and $n=1$.}, we revisit this problem.  In this {\\it Letter}, we derive new constraints from the observed properties of known high-redshift quasars, and assess the potential for tightening these constraints with future quasar surveys. ",
        "conclusions": "Existing data on high-redshift quasars implies that BHs as massive as $\\sim 3\\times10^9~{\\rm M_\\odot}$ were assembled when the universe was only a tenth of its present age.  Figure 2 shows that the massive BH inferred for SDSS 1044-0125 at $z=5.8$ can grow in hierarchical galaxy formation models with plausible parameter values for the initial seed mass ($\\sim 10~{\\rm M_\\odot}$), the minimum velocity dispersion of collapsed objects that harbor such a seed ($\\ga 50~{\\rm km~s^{-1}}$), the radiative efficiency ($\\sim 6\\%$), and the luminosity in Eddington units ($\\sim 1$). Figure 3 illustrates the upper envelope of plausible luminosity values for high-redshift quasars that might be found in future surveys.  There are several caveats to the constraints we derived.  First, beaming or lensing may affect the apparent magnitude of some of the brightest and highest redshift quasars (Barkana \\& Loeb 2000b).  Second, we have assumed that $\\eta$ and $\\epsilon$ maintain the same values during the luminous quasar phase and the main growth phase of the BH mass. It is possible, however, that $\\epsilon$ is negligibly small during the early growth phase of the BH, and high during the luminous quasar phase (e.g. Haehnelt \\& Rees 1993).  This would be equivalent to a corresponding increase in the seed BH mass in our model. Multiple seeds per halo which eventually coalesce are also equivalent to a single massive seed with the sum of their masses.  Finally, although a recent cosmological simulation finds good agreement with the Press--Schechter mass function for the lowest halo masses of interest here at $z=10$ (Jang-Condell \\& Hernquist 2000); the accuracy of this ansatz still remains to be tested over a wider range of redshifts and halo masses, where there might be systematic deviations (Jenkins et al. 2000; Sheth \\& Tormen 2000).  Recently, various modifications to the standard $\\Lambda$CDM model have been proposed because of a potential conflict between theory and observations on small spatial scales (e.g., Kamionkowski \\& Liddle 2000; Bode, Ostriker \\& Turok 2000; Barkana, Haiman \\& Ostriker 2000, and references therein).  Our constraints are expected to tighten significantly in such models, since they generically reduce small scale power and thus eliminate potential hosts of seed BHs at high redshifts. Future surveys of high-redshift quasars will test these models."
    },
    "0011/astro-ph0011373_arXiv.txt": {
        "abstract": "We present measurements of the optical/UV emission lines for a large homogeneous sample of 993 quasars from the Large Bright Quasar Survey. Our largely automated technique accounts for continuum breaks and galactic reddening, and we perform multicomponent fits to emission line profiles, including the effects of blended iron emission, and of absorption lines both galactic and intrinsic.  Here we describe the fitting algorithm and present the results of line fits to the LBQS sample, including upper limits to line equivalent widths when warranted.  The distribution of measured line parameters, principally \\ew\\, and FWHM, are detailed for a variety of lines, including upper limits.  We thus initiate a large-scale investigation of correlations between the high energy continuum and emission lines in quasars, to be extended to complementary samples using similar techniques. High quality, reproducible measurements of emission lines for uniformly selected samples will advance our understanding of active galaxies, especially in a new era of large surveys selected by a variety of complementary methods. ",
        "introduction": "\\label{intro}  Quasars, some of the most luminous objects in the Universe, allow us to observe towards the start of cosmic time, probing extremes of physics, and illuminating  the intervening matter.  Though most of the QSOs now known were discovered by their unique spectral energy distributions (SEDs) or emission lines, several interrelated and fundamental questions remain:  \\smallskip  1)  What significant correlations exist between the observed SED and emission lines in large uniform quasar samples?  2)  How are the observed SED and emission lines affected by intrinsic absorption and orientation?  3)  Can the shape of the intrinsic SED be determined?  4)  What effect does the intrinsic SED have on emission lines and what does that tell us about physical conditions in the QSO?  \\smallskip  The production of emission lines in QSO spectra is widely attributed to photoionization and heating of the emitting gas by the UV to X-ray continuum (e.g., Ferland \\& Shields 1985; Krolik \\& Kallman 1988). Emission lines from a given ion are particularly sensitive to photons of energy above the corresponding ionization threshold, but ionization from excited states and heating via free-free and H$^-$ absorption also determine the line's principal ionizing/heating continuum. Many important lines respond to the extreme ultraviolet (EUV) or soft X-ray continuum.  Unfortunately, the EUV band is severely obscured by Galactic absorption, although some constraints on the EUV ionizing continuum are available through analysis of the adjacent UV and soft X-ray windows.  The overall similarity of QSO emission line spectra has on occasion been taken as evidence of fairly uniform physical conditions in the broad emission line region (BELR).  At first, this encouraged the assumption that clouds in the BELR inhabit a narrow swath of parameter space in density, size, and ionization parameter ($U$).  Observed correlations of line equivalent widths (hereafter, \\ew) with SEDs present challenges to geometric and photoionization models that need to be answered.  Photoionization pioneers such as Mushotzky \\& Ferland (1984) ran models on a single cloud, while photoionization models assuming a distribution of clouds (i.e., not a single $U$) predict a strong dependence of $W_{\\lambda}$ on the (assumed power law) slope of the SED (Binette et al. 1989; Baldwin et al. 1995; Korista, Baldwin, \\& Ferland 1998).  Large uniform samples of QSO emission line measurements are key to test photoionization models, but are also necessary for other important tests of quasar triggering and evolution.  Tests of the binary vs. lens hypothesis for quasar pairs at the same redshift require quantitative measures of the probability of finding the measured spectral similarity derived from complete samples (Mortlock, Webster \\& Francis 1999).  Since optically-selected QSOs at similar redshift may tend to be more similar than 2 QSOs chosen entirely at random, studies of QSO spectral properties should be based on samples large enough to provide statistically significant subsamples across a wide range of luminosity and redshift (Peng et al. 1999).  In a small, complete sample of optically-selected QSOs studied by Laor et al. (1997), the strongest correlation found between X-ray continuum and optical emission line parameters was of the soft X-ray spectral slope (\\ax), and the FWHM of H$\\beta$.  Strong correlations between \\ax~ and the relative strengths of [\\ion{O}{3}] $\\lambda$5007 and iron emission were seen both there and in previous studies, particularly Boroson \\& Green (1992) (hereafter BG92). A principal component analysis (PCA) was employed by BG92 to examine an unwieldy set of emission line and continuum correlations within a sample of 87 QSOs taken from the Palomar-Green Bright Quasar Survey (BQS; Schmidt \\& Green 1983). This indicated that the strongest sources of variance in the optical spectra, the ``primary eigenvector'', involved an anti-correlation between the strength of [\\ion{O}{3}] $\\lambda$5007 and iron emission that was also linked to the FWHM of H$\\beta$. This work was extended by Marziani et al. (1996) (hereafter M96) and Sulentic et al. (2000) to include a study of the spectral region around \\ion{C}{4} $\\lambda$1549. These papers identified Narrow-line type 1 Seyfert galaxies (NLS1) and steep spectrum radio galaxies as occupying opposite ends of the primary eigenvector.  Other work utilizing the high spectral resolution of the {\\em Hubble Space Telescope}~ Faint Object Spectrograph ({\\em HST} FOS) showed a link between \\ion{Al}{3} $\\lambda$1859, \\ion{Si}{3}] $\\lambda$1892, and \\ion{C}{3}] $\\lambda$1909 emission line strengths and the primary eigenvector but suffered from the small sample size and a focus on NLS1 or the low redshift radio quiet objects in the BQS (Aoki \\& Yoshida 1999; Kuraszkiewicz et al. 2000).  The samples of AGN used by most authors are almost entirely taken from BG92 and so do not represent a statistically complete sample and may not be truly representative of the QSO population (Wisotzki et al. 2000). They also deal with the strongest emission features in the OUV spectrum and, while several hypotheses have been proposed, the physical origin of the observed correlations remains a puzzle.  There is also a concern that biases within these small samples may have exaggerated the significance of the primary eigenvector.  Even in the clear presence of extrinsic effects such as absorption along the line of sight, emission lines have been used to infer the strength and shape of the high energy SED.  As an example, the similarity of UV emission-line properties in broad absorption line (BAL) and non-BAL QSOs (Weymann et al. 1991) has been cited as evidence that orientation is the cause of the BAL phenomenon (i.e., that {\\em all} radio-quiet QSOs have BAL clouds).  However, BALQSOs are now known to exhibit markedly weak X-ray emission as a class (Green \\& Mathur 1996; Gallagher et al. 1999; Brandt, Laor, \\& Wills 2000).  But if similar emission lines indeed vouch for similar intrinsic high energy SEDs, then the large X-ray to optical continuum flux ratio (\\aox) values observed for BALQSOs are likely to be caused by strong absorption along the line-of-sight rather than by differences in their intrinsic SEDs.  In fact, the soft X-ray deficit requires absorption at least an order of magnitude larger than those estimated from optical/UV (OUV) spectra alone and the UV and X-ray absorbers have yet to be definitively identified as one (e.g., see discussion in Mathur, Wilkes, \\& Elvis 1998).  Since BALQSOs may be heavily absorbed, the question of whether similar emission lines are testimony for similar intrinsic SEDs must be answered through the systematic study of line and continuum correlations in {\\em unabsorbed} QSOs.  Even in supposedly unabsorbed QSOs, a strong inverse correlation exists between optical iron emission, and \\ion{O}{3} $\\lambda$ 5007 in QSOs (BG92).  Iron emission is under-predicted by a large factor in standard photoionization models relative to observations (Collin \\& Joly 2000), and the continuum source responsible for creating rest-frame optical iron emission lines comes almost entirely from X-rays above about 1keV.  QSOs with large \\aox\\, of which BALQSOs represent an extreme, have not only weak \\ion{O}{3}, but feeble narrow line emission in general (Green 1996).  We suggest that the observed correlations may be an effect of absorption of the intrinsic high energy continuum by thick clouds interior to the NELR.  The X-ray heating of these clouds makes them emit copious \\ion{Fe}{2}, lines, but prevents the ionizing continuum from reaching the NELR.  This is reinforced by another observed correlation - that the most highly absorbed BALQSOs, the low-ionization BALQSOs, albeit X-ray quiet, are generally very strong optical \\ion{Fe}{2}, emitters (Stocke et al. 1992; Lipari, Terlevich, \\& Macchetto 1993). Our interpretation needs to be tested across a range of absorption, in large representative QSO samples.  We have therefore undertaken a major study of quasar line emission, with careful accounting for absorption lines and blended iron emission, using largely automated procedures on carefully-selected samples. The analysis of samples of QSOs with more than a few hundred spectra requires some amount of automation to give consistent results. The largest sample of QSO spectra currently available is that of the Large Bright Quasar Survey and here we describe the initial results from the measurement of this sample and the analysis methods that will also be applied to other large samples of QSO spectra that will become available in the near future. ",
        "conclusions": "To examine the robustness of the analysis procedure in OUV spectra with improved S/N and spectral resolution we have modeled two composite QSO spectra using the techniques described above. The first was constructed by Francis et al. 1991 (hereafter F91) from 718 spectra in the LBQS sample itself, and the second from {\\em HST} FOS observations of 101 QSOs presented in Zheng et al. (1997) (hereafter Z97). We find that even where the limit of two Gaussian components for an emission line does not reproduce the line profile very well, the measured \\ew~ is within 10\\% of values tabulated in the above papers. The exception is for \\ion{Mg}{2} where the F91 measurement is higher and the Z97 measurement is lower than that measured using the techniques presented above. Both discrepancies can be attributed to the presence of iron emission features blending into the wings of the \\ion{Mg}{2} emission line. Although F91 fitted a global continuum and corrected for iron emission using local spline fits to emission complexes, the \\ew~ was measured as the total flux above the continuum and iron emission across more than 50\\AA~ in the spectrum and so may include some hidden iron emission. This is also likely to be the explanation of a similar discrepancy in the measurements of H$\\gamma$ + [\\ion{O}{3}] $\\lambda$4363 and [\\ion{O}{3}] $\\lambda$4959 in the F91 composite.  The measurements in Z97 use a local continuum which is difficult to determine accurately in regions where iron emission complexes are strong, and may result in their lower \\ew~ measurement. This confirms that correct accounting for iron emission is crucial in the measurement of emission lines in QSO spectra.  We can also compare the global properties of the LBQS emission line measurements with the measurements of these composite QSO spectra. The Z97 composite shows larger \\ew~ in Lyman $\\beta$ + \\ion{O}{6}, Lyman $\\alpha$, and \\ion{C}{4}.  The discrepancy remains even if the means of only our detected emission lines are compared. (The inclusion of the censored data will naturally lower the means of a distribution, but has little effect on the distribution of strong lines like Lyman $\\alpha$ and \\ion{C}{4} which are well constrained in nearly all cases).  We suspect that the difference is attributable to the large (60\\%) fraction of radio loud objects in the heterogeneous sample used to construct the Z97 composite. The LBQS sample is predominantly Radio Quiet (RQ) with $\\sim$ 10\\% Radio Loud (RL) (Hooper et al. 1995) and Z97 showed that there are significant differences in the strength of these lines between RL and RQ populations, independent of luminosity. A marginally significant RL/RQ difference was also seen in a sample of 255 of the optically brightest QSOs in the LBQS sample (Francis, Hooper, \\& Impey 1993). A more detailed examination of this effect will be made in a subsequent paper that will contrast the continuum and emission line properties within the full LBQS sample.  In contrast, the \\ew~ of the emission features in the F91 composite are lower than the mean and medians of the measured emission lines in the LBQS sample.  Although the F91 composite was created using a large proportion of the same spectra measured here, it is not surprising that the measurements should differ. Lyman\\,$\\alpha$, which shows the largest dispersion of relative flux within the sample (see Fig. 5 in F91), is most affected by narrow absorption features.  Allowance for the effects of narrow absorption increases the \\ew~ modeled with Gaussian components relative to an estimate based on the integration of flux above a continuum. For [\\ion{Ne}{3}] and H$\\delta$, the F91 composite indicates a stronger feature than suggested by the KM mean presented in Table~7. However, the measurements are similar when only the detected lines are studied.  There has been no investigation of inclusion of censored data on the emission line and continuum correlations, particularly the Baldwin effect (Baldwin 1977; Zamorani et al. 1992; Osmer, Porter, \\& Green 1994) and the correlations associated with the primary eigenvector of spectral variance in AGN (BG92, M96, Sulentic et al. 2000).  Previous studies have focussed on a small number of bright emission features; particularly \\ion{C}{4} $\\lambda$1549, \\ion{He}{2} $\\lambda$4686, H$\\beta$, \\ion{O}{3} $\\lambda$5007 and iron emission.  We briefly summarize the published measurements for QSO samples of BG92 (87 low z QSOs, optical spectra), McIntosh et al. (1999) (32 high z QSOs with IR spectra), M96 (52 QSOs with optical and UV spectra) and Sulentic et al. 2000 (125 QSOs) in Table~8. Where upper limits are quoted, we have used survival analysis to enable a more direct comparison to our results in Table~7. Some of the published work has also been presented separately as RL or RQ subsamples.  While all these studies suffer from differing selection effects, the comparison samples also differ from the LBQS sample in redshift range and magnitude limits.  The most notable difference between the properties of the LBQS and the measurements shown in Table~8 can be seen in the relative strength of optical iron emission, the mean value of which appears significantly higher in BG92 and MS96.  The region used to measure the \\ew~ of optical iron emission is identical to that used in BG92 (and subsequent authors) and a test using the methods described in \\S3 to measure the BG92 sample spectra (kindly supplied by T. Boroson) shows reasonable ($\\sim$ 20\\%) agreement in \\ion{Fe}{2} \\ew.  The measurement of UV iron emission in M96 appears similar to the mean of the distribution of iron emission strength measured in the LBQS sample. However the window used by M96 to characterize the iron emission strength was 1550 -- 1750\\AA~ and the iron template flux in this region is approximately 6.5\\% of the flux in the region used for the study of the LBQS sample presented here (i.e. the mean UV iron emission in M96 is 15 times higher than the mean in the LBQS sample!). The unusually strong iron emission was noted by M96 as a possible bias in their sample.  The heterogeneous sample (52 AGN for which $HST$ FOS spectra covering \\ion{C}{4} were available in 1994 and for which the authors had matching optical spectra) has an overabundance of strong iron emitters relative to the BG92 sample. The difference in iron emission distributions contrasts to the similarity seen in the distributions of \\ew~ for [\\ion{O}{3}]$\\lambda$5007. The strongest correlations reported for optical spectra of AGN involve iron emission, H$\\beta$ and [\\ion{O}{3}].  We expect that our new, more homogeneous emission line measurements, of iron in particular, may significantly reshape the discussion of correlations that produce the largest variance in PCA analysis of quasar emission lines."
    },
    "0011/astro-ph0011145_arXiv.txt": {
        "abstract": "We present observations of the 1997 outburst of the X-ray transient GS 1354$-$64 (BW Cir) at X-ray, optical and, for the first time, radio wavelengths; this includes upper limits to the linear and circular polarisation of the radio data. The X-ray outburst was unusual in that the source remained in the low/hard X-ray state throughout; the X-ray peak was also {\\em preceded} by at least one optical outburst, suggesting that it was an `outside-in' outburst -- similar to those observed in dwarf novae systems, although possibly taking place on a viscous timescale in this case. It therefore indicates that the optical emission was {\\em not} dominated by the reprocessing of X-rays but that instead we see the instability directly. While the radio source was too faint to detect any extended structure, spectral analysis of the radio data and a comparison with other similar systems suggests that mass ejections, probably in the form of a jet, took place and that the emitted synchrotron spectrum may have extended as far as infrared wavelengths. Finally, we compare this 1997 outburst of GS 1354$-$64 with possible previous outbursts and also with other hard state objects, both transient and persistent. It appears that a set of characteristics -- such as a weak, flat spectrum radio jet, a mHz QPO increasing in frequency, a surprisingly high optical:X-ray luminosity ratio and the observed optical peak preceding the X-ray peak -- may be common to all hard state X-ray transients. ",
        "introduction": "Soft X-ray transients are a class of low mass X-ray binaries in which instabilities in the accretion disc cause a sudden increase in mass accretion rate onto the compact object, resulting in an outburst (see e.g. Charles 1998 and references within). While reminiscent of the dwarf novae class of cataclysmic variables (Van Paradijs 1996) these outbursts occur with intervals typically 10--20 years and usually decline over several months (Tanaka \\& Lewin 1995; Van Paradijs \\& McClintock 1995), sometimes displaying secondary outbursts as they do so (e.g. the X-ray flux of GS 1124$-$684 rose again by a factor of $\\sim$ 2 about 80 days after the initial outburst. See Tanaka \\& Shibazaki 1996 and references within). In between outbursts the X-ray transients lie in a quiescent state and this gives us an excellent opportunity to study the companion star.  Although these objects are frequently referred to as {\\em soft} X-ray transients on account of their ultrasoft X-ray spectra during the outbursts, an increasing number of their outbursts do {\\em not} show the soft component (e.g. GRO J1719$-$24, Van der Hooft et al. 1996). Instead the power law component dominates the energy spectrum, reminiscent of the low/hard X-ray state of the persistent source Cyg X-1. However, as the term `hard X-ray transient' has previously been used to describe the Be star+neutron star binaries (which typically show hard X-ray spectra) we do not use this term here. It is important to note that the term `low/hard state' is used for historical reasons, dating back to observations of Cyg X-1 during the 1970's -- it is somewhat misleading as its current definition (Van der Klis 1995) is based on spectral properties and the source flux need not necessarily be `low'.  \\begin{figure*} \\vspace*{-1cm} \\begin{center} \\leavevmode \\psfig{file=lightcurve.eps, width=17cm,angle=0} \\vspace*{-.5cm} \\caption{BATSE, ASM, PCA, optical and radio (4.8 GHz) lightcurves of the 1997 outburst for the period 8 October 1997 -- 16 May 1998. Radio lightcurves at other frequencies are shown later in this paper. Error-bars on the optical points are smaller than the symbols. The variability of the QPO frequency is also shown in the bottom panel.} \\label{lightcurves} \\end{center} \\end{figure*}   A large number of the X-ray transients have been classified as black hole candidate (BHC) X-ray binaries -- where possible, the preferred method is to determine the mass function (e.g. Charles 1998 and references within) but in a large number of cases a black hole nature has been suggested on account of the X-ray properties. Following the unification scheme for the X-ray power density spectra of neutron stars and BHCs proposed by Van der Klis (1994; 1995), the power spectra of BHCs can be sorted according to the mass accretion rate; this is thought to increase from the low/hard state, through the intermediate state and the high/soft state, up to the very high state (but see also Homan et al. 2000). During the low/hard state strong band-limited noise is present. In this scheme low frequency ($\\leq$ 1 Hz) quasi-periodic oscillations (QPOs) seen in the low/hard state are seen as an aspect of the low/hard state noise. Recently, Psaltis, Belloni, \\& Van der Klis (1999) reported correlations between the frequencies of QPOs and broad band power spectral features of neutron star and black hole X-ray binaries.  Transitions between X-ray states have been observed during the rise and subsequently during the decay of an outburst for a number of sources (e.g. GS 1124$-$68 (Ebisawa et al. 1994), GRO J1655$-$40 (M\\'endez et al. 1998)). Of particular interest is GX 339$-$4, in which the transition from low/hard to high/soft X-ray states and back again was clearly tracked at radio wavelengths, with radio emission strongly suppressed in the high/soft state (Fender et al. 1999; Fender 2000a).  Ground based follow up observations of X-ray transients have shown that outbursts also occur at optical and radio frequencies. A0620$-$00, GRO J1655$-$40, GS 2023+338 (V404 Cyg) and GS 1124$-$684 are examples of X-ray transients detected in all three wavelength regimes (Tanaka \\& Lewin 1995, and references within); in particular, the X-ray, optical and radio behaviour for GS 2023+338 were correlated during the 1989 outburst (Han \\& Hjellming 1992) and it is interesting to note that during this outburst (as in the case of GS 1354$-$64) the source remained in the low/hard state (e.g. Tanaka \\& Lewin 1995).  \\begin{table*} \\caption{Log of the $RXTE$/PCA observations. The penultimate column indicates the category of the observation (see text).} \\label{obs_log} \\begin{tabular}{ccccccc} \\hline Number & Observation & Date \\& & Total on source & Category & Mean count\\\\ &    ID       &  & observing time (ksec.) & & rate (cts/s) \\\\ \\hline 1 & 20431-01-01-00       & 1997-11-18 & 3.2 & rise & $\\sim$750\\\\ 2 & 20431-01-02-00       & 1997-11-19 & 6.0 & rise & $\\sim$875\\\\ 3 & 20431-01-03-00       & 1997-11-22 & 8.4 & rise & $\\sim$900\\\\ 4 & 20431-01-04-00       & 1997-12-05 & 7.6 & top & $\\sim$950\\\\ 5 & 20431-01-05-00       & 1997-12-12 & 6.4 & top & $\\sim$860\\\\ 6 & 30401-01-01-00       & 1997-12-27 & 7.0 & decay1 & $\\sim$700\\\\ 7 & 30401-01-02-00       & 1998-01-04 & 7.7 & decay2 & $\\sim$575\\\\ 8 & 30401-01-03-00       & 1998-01-12 & 6.1 & decay2 & $\\sim$550\\\\ 9 & 30401-01-04-00       & 1998-01-26 & 7.1 & decay3 & $\\sim$425\\\\ \\hline \\end{tabular} \\end{table*}  The optical emission is thought to be the result of reprocessing of X-rays in the disc -- the disc absorbs X-rays and re-emits them at lower energies, delaying the optical by a few seconds as reported by Hynes (1998) in the case of GRO J1655$-$40. The radio appears to be the result of beamed synchrotron emission taking place in relativistic ejections of material from the centre of the disc; this may be in response to additional material being transported through the disc by the instability and so we might expect some correlation between the radio and the X-ray/optical behaviour. However, a comparison between the radio lightcurves of A0620$-$00, GS 1124$-$68 and GS 2000+25 suggests that their profiles depend on inclination angle, the number of ejections, the time between them and their strength and speed (Kuulkers et al. 1999). Therefore, as the radio emission may be anisotropic in the observer's frame, unlike (presumably) X-ray and optical, we might instead expect that it is {\\em not} observed to be so closely correlated with these two components as they are with each other.  \\subsection{GS 1354$-$64}  On 1987 February 13, the ASM on board the {\\it Ginga} satellite (Swinbanks 1987) discovered GS 1354$-$64, observing it until 1987 August (Makino et al. 1987). The X-ray data were well fit by a soft disc blackbody component, with inner disc temperature of $\\sim$0.7 keV, and a hard power law with photon index 2.1 (Kitamoto et al. 1990). These X-ray properties are common for X-ray transients during outburst (Tanaka and Lewin 1995) and hint at a black hole nature for the compact object. It was also observed in the optical during the outburst, yielding magnitudes of $V\\sim16.9$, $B\\sim18$ (Kitamoto et al. 1990). GS 1354$-$64 had not previously been observed in the radio.  \\begin{table*} \\caption{Identifiers, coordinates (J2000 equinox) and magnitudes of the eight reference stars used in the photometry. The `BJF' part of the identifier refers to Brocksopp, Jonker, Fender et al. 2000 (this work)} \\label{refstars} \\center \\begin{tabular}{cccc} \\hline Identifier, RA, Dec.&$V$ (mags)&$B-V$ (mags)&$V-R$ (mags) \\\\ \\hline BJF J135807.07$-$644309.9&16.269 (0.005)&   0.987 (0.016)&   0.632 (0.006)\\\\ BJF J135807.94$-$644309.3&16.028 (0.004)&   1.111 (0.013)&   0.646 (0.005)\\\\ BJF J135815.53$-$644305.1&15.446 (0.002)&   1.905 (0.014)&   1.088 (0.002)\\\\ BJF J135814.93$-$644337.6&15.797 (0.003)&   1.570 (0.015)&   0.895 (0.003)\\\\ BJF J135813.21$-$644348.4&15.506 (0.002)&   1.914 (0.015)&   1.084 (0.002)\\\\ BJF J135809.39$-$644439.1&15.463 (0.003)&   0.889 (0.006)&   0.557 (0.003)\\\\ BJF J135807.71$-$644513.1&15.036 (0.002)&   0.937 (0.005)&   0.559 (0.002)\\\\ BJF J135811.48$-$644516.0&15.026 (0.002)&   0.811 (0.005)&   0.485 (0.002)\\\\ \\hline \\end{tabular} \\end{table*}  The position of this source on the sky is consistent with the position of two other transient sources; Cen X-2 (Francey et al. 1971) and MX 1353$-$64 (Markert et al. 1979). Cen X-2 was the first `soft X-ray transient' discovered and one of the brightest (Tanaka \\& Lewin 1995). The outbursts of MX 1353$-$64 (1972) and GS 1354$-$64 (1987 and 1997) reached much lower intensities and showed different X-ray spectral properties -- if Cen X-2 and MX 1353$-$64 were the same source as GS 1354$-$64 then it must show at least four different states (Kitamoto et al. 1990). This is not unfeasible -- e.g. GX 339$-$4 has been observed in four states, Cyg X-1 in three (see Fender 2000a, and references within) -- it is therefore probable that these sources are indeed one and the same, in which case we note that the 1997 outburst was significantly sub-Eddington. The error circles for the positions of each `source' can be found in Kitamoto et al. (1990). We investigate the nature of the various states further in Section 4.  GS 1354$-$64 was observed in outburst again in 1997 November. Preliminary results from the ASM on board the Rossi X-ray Timing Explorer ({\\it RXTE}) showed a rise from 16 to 50 mCrab in the first half of November (Remillard, Marshall \\& Takeshima 1997); the BATSE instrument on the Compton Gamma Ray Observatory ({\\it CGRO}) detected hard X-rays up to 200 keV which rose from 60 to 160 mCrab (Harmon \\& Robinson 1997). The optical and infrared counterparts were also detected with magnitudes of $R\\sim16.9$, $B\\sim18.1$ (Nov. 22; Castro-Tirado, Ilovaisky \\& Peterson 1997) and $J\\sim15.35$, $K\\sim13.95$ (Nov. 20; Soria, Bessell \\& Wood 1997). Optical spectroscopy during the decay revealed strong emission lines corresponding to H$\\alpha$, H$\\beta$ and He\\,{\\sc ii}\\, $\\lambda$4686 plus weaker He\\,{\\sc i}\\,$\\lambda$4471 and H$\\gamma$ emission (Buxton et al. 1998). The H$\\alpha$ emission profile varied from single- to double-peaked over three nights of observations during the decay from outburst, indicating the presence of an accretion disc. The radio counterpart was detected at 2.5, 4.8 and 8.6 GHz with a flux of 1--3 mJy (Nov. 24/25; Fender et al. 1997). No indication of the orbital period was detected at any wavelength.  A recent paper by Revnivtsev et al. (2000) gives an analysis of $RXTE$/ASM, HEXTE and PCA data during the outburst. Their work indicates that the X-ray energy spectrum is dominated by a power law, typical of X-ray binaries in the low/hard state. However, the fit is improved by the addition of a 6.4 keV iron line and a reflection component suggestive of the presence of an accretion disc; low energy disc photons are upscattered by a Comptonising corona to higher energies, resulting in the X-ray emission. Here we present $B$, $V$ and $R$ photometry and radio data from the outburst, combining it with further $RXTE$/PCA analysis and public ASM and $CGRO$/BATSE data. ",
        "conclusions": "Our multiwavelength dataset has enabled us to make a number of suggestions as to what took place during the outburst of GS 1354$-$64 in 1997 November. The optical and radio lightcurves suggest that there were two maxima during the outburst, the first prior to our observations and the second at $\\sim$JD2450784, depending on wavelength. There is also the hint of a third optical rise just following the X-ray peak (JD2450796) -- by this time it is possible that the disc is sufficiently bright in X-rays for the third optical brightening to be the result of X-ray reprocessing in the disc (although this is not confirmed by the hard X-ray energy spectrum). For future outbursts, simultaneous X-ray and high time resolution optical observations would be extremely beneficial; discovery of the mHz QPO (mentioned in Section 3.1) in the optical would help to determine the efficiency of the reprocessing of soft X-rays.  It is clear that the optical lightcurve reached a local maximum {\\em before} the X-rays peaked. While common in dwarf novae outbursts, a preceding optical peak is not generally seen in X-ray transient events -- quite possibly only due to the fact that it is the X-ray satellites which tend to discover them. If this is the case then it is unlikely that reprocessing of X-rays was the dominant source of optical emission. (However, it is possible if there is some non-straightforward disc geometry e.g. GRO J1655$-$40 (Hynes et al. 1998)). Our comparison of `soft' X-ray transients which remain in the low/hard X-ray state for the duration of the outburst shows that GS 2023+338 and possibly GRO J0422+32 were also observed to peak in the optical before the X-ray. While this may be coincidence in timing of the observations, we consider the possibility that in these hard state transients the X-ray emission is at a sufficiently low level for the inwards travelling instability to dominate the optical emission; in outbursts with a soft X-ray component reprocessing of X-rays would dominate instead. This is confirmed by the surprisingly high optical/X-ray luminosity ratios that these low/hard state sources have -- in a `normal' soft event the relative X-ray luminosity would be much higher than in the hard state sources.  If the peak of the X-ray lightcurve corresponds to the time at which the instability reaches the hot inner regions of the disc then a disc-crossing time of $\\sim$10 days is inferred. For this to be true then the disc and therefore the orbit must be large. The fact that no orbital period has been determined for this source also suggests that the orbit may be large (this may also contribute to the possible over-bright optical emission mentioned in Section 3.2) -- although it is more likely that there is insufficient optical data for any periodicities to be determined. If indeed the orbit {\\em is} large then BW Cir would have to be an evolved star filling its Roche Lobe, suggesting a similar type of system to GRO J1655$-$40.  This inferred disc-crossing time is very long and it is unlikely that the disc of a system with such a low soft X-ray luminosity and hard X-ray spectrum would extend to a small inner radius. It is also unlikely that the instability could trigger a radio ejection before reaching the innermost regions of the disc and causing the X-ray peak. The increase in QPO frequency suggests that the increase in soft X-ray luminosity is the result of the inner disc radius decreasing; the surface area of the disc and the temperature of the inner regions are therefore increasing. This would be consistent with other X-ray binaries -- the low/hard state is characterised by a large inner disc radius and this decreases on transition to the high/soft state (e.g. Esin et al., 1998).  Therefore it appears that the instability moves inwards, emitting in the optical regime. As the transported mass fills the `hole' at the centre of the disc on a viscous timescale, the soft X-ray emission from the corona increases. This scenario is comparable with that suggested by Shahbaz et al. (1998) to explain the optical/X-ray delay of the 1997 outburst of the neutron star X-ray transient Aql X-1.  If we assume that a radio event is triggered when the instability reaches the inner edge of the disc then this considerably reduces the required disc crossing time (i.e. the disc crossing time becomes the delay between the optical and radio peaks rather than the delay between the optical and X-ray) and seems more likely. From observations of other sources in the hard state it is probable that the first radio event is a discrete ejection, with relativistic bulk velocity, which rapidly evolves to an optically thin spectrum. Subsequently, in the low/hard state, a powerful, steady, partially self-absorbed jet is produced which is observed approximately as a flat-spectrum component (Fender 2000b). It is also possible that repeated discrete ejections take place -- however, these are more commonly seen in the soft X-ray transient outbursts and do not seem to be a common feature of the low/hard X-ray state.  Alternatively (or additionally?), we also note that the optical and radio data have comparable fluxes. It is therefore possible that an ejection event took place, emitting synchrotron radiation at optical and increasingly longer wavelengths as the ejected material became optically thin at increasing distances from the central source. The flat spectrum (Fig.~\\ref{spectrum}), apparently extending from radio wavelengths to the $R$ band, would certainly support this scenario. If, as our value of $\\xi$ suggests, the optical component is over-bright then synchrotron should not be discounted as a possible contributor to the optical emission in addition to the thermal disc spectrum.  It is important to consider the implications of a jet (or mass ejection) that emits flat spectrum synchrotron radiation over such a wide range of frequencies. While there are a number of models that work reasonably well for the `flat spectrum' AGN, none of these is consistent with the {\\em much flatter} radio -- (sub-)mm spectra of e.g. Cyg X-1, GRS 1915+105 and Cyg X-3 (Fender et al. 2000 and references within). The most convincing explanation for the flat spectrum is synchrotron radiation from a partially self-absorbed jet. With the power of the jet directly proportional to the bandwidth, detection of a high frequency cut-off to the flat spectrum would be very useful in determining the power of the jet -- this has not been found in any of the flat spectrum X-ray binaries.  In summary, our limited observations suggest that the 1997 outburst of GS 1354$-$644 was the result of an instability crossing the disc on a slow viscous timescale, emitting at optical wavelengths. As the instability reached the (large) inner radius of the disc it triggered a mass ejection which emitted synchrotron radiation at radio wavelengths and possibly through to the near infrared (and optical?) regime. The X-ray source remained in the low/hard state throughout, gradually becoming brighter as the instabilities carried more matter into the hot inner regions and subsequently decreasing the inner disc radius. These results clearly reflect the importance of simultaneous optical, infrared and radio studies of future X-ray transient events."
    },
    "0011/astro-ph0011309_arXiv.txt": {
        "abstract": "We study the sensitivities of space infrared interferometers. We formulate the signal-to-noise ratios of infrared images obtained by aperture synthesis in the presence of source shot noise, background shot noise and detector read noise. We consider the case in which $n$ beams are pairwise combined at $n(n-1)/2$ detectors, and the case in which all the $n$ beams are combined at a single detector. We apply the results to future missions, Terrestrial Planet Finder and Darwin. We also discuss the potential of a far-infrared interferometer for a deep galaxy survey. ",
        "introduction": "Early in the 21st century space infrared observatories such as ASTRO-F\\cite{astro-f}, Space Infrared Telescope Facility (SIRTF)\\cite{sirtf}, Far Infrared and Submillimeter Telescope (FIRST)\\cite{first}, and Space Infrared Telescope for Cosmology and Astrophysics (SPICA)\\cite{nak98} are planned to study the formation of planets, stars, and galaxies. The improvement in sensitivity is so great that the source detection limits of these space observatories are expected to be set by source confusion, especially in the far infrared. To further improve the source detection limits, individual sources need to be resolved by an infrared interferometer. Therefore it is important to quantify the sensitivity of a space infrared interferometer for the purpose of synthesis imaging.  Following the observatories, space infrared interferometers such as the Terrestrial Planet Finder (TPF)\\cite{tpf} and Darwin\\cite{darwin} are considered,  with an emphasis on the detection of terrestrial planets around nearby stars. TPF and Darwin will consist of several radiation cooled apertures and will be capable of synthesis imaging as well as planet detection by nulling interferometry. The detection limits of these interferometers for general synthesis imaging are of great interest. Among many applications of  synthesis imaging with a space infrared interferometer, we are  particularly interested in deep galaxy count by which the history of galaxy formation can be studied. We discuss how a TPF-like interferometer operated in the far infrared can improve the number count data of high redshift galaxies.   The beam combination geometry is a major issue in studying the sensitivity of an interferometer. One extreme is the $^nC_2$ interferometer, in which $n$ beams are divided into $n(n-1)$ subbeams and they are pairwise combined at $^nC_2 = n(n-1)/2$ detectors. There is one detector for each baseline. The other extreme is the $^nC_n$ interferometer in which all the $n$ beams are combined at $^nC_n = 1$ detector. Prasad and Kulkarni\\cite{pra89} (henceforth PK89) have rigorously analyzed the SNRs of the $^nC_2$ and $^nC_n$ interferometers for spectrally narrow  light in the presence of source shot noise  and concluded that the sensitivities of both interferometers are essentially the same apart from a numerical factor. The results by PK89 are applicable to optical interferometry, but not to infrared interferometry because of the presence of thermal background and detector read noise. The beam combination geometry can have an impact on the sensitivity in the infrared for instance because of the number of detector pixels which affects the total read noise. The beam combination geometry may also have a major impact on the architecture of the space interferometer.    Baseline redundancy is a factor that affects the SNR of a synthesis image for a given number of apertures. Since the number of apertures deployed or flown separately in space is likely to be limited, it is important to gain $uv$ coverage by using nonredundant baselines. Throughout this paper, we only consider interferometers with nonredundant baselines.  There are two types of noise which appear in aperture synthesis, one of which is removable while the other is intrinsic and not removable. Sampling noise is potentially introduced by the incompleteness of $uv$ coverage. However, deconvolution techniques such as CLEAN and Maximum Entropy Method (MEM)\\cite{per89}, when applied to radio interferometric data, appear to compensate for sampling noise. We thus regard sampling noise as removable noise. On the other hand, noise resulting from the photodetection process is not removable. The SNR of a dirty image to which deconvolution has not been applied is limited by photodetection noise. Here we derive the SNR of a dirty image as determined by source shot noise, thermal background shot noise, and detector read noise, and equate it to the SNR of the synthesized image.  There are two different methods by which a synthesized image can be constructed from the visibility data: inversion without total counts and true inversion. In the former, the zero frequency phasor is neglected, implying that the total number of photons in the synthesized image is zero. Despite this seemingly unattractive feature, this is the standard method in radio astronomy. In the latter method, the van Cittert-Zernike theorem is strictly applied and the zero frequency phasor as well as the $n(n-1)/2$ complex phasors is used. The image produced by this technique has the desirable property of nonnegativity. In the presence of additive thermal background which overwhelms the signal, the magnitude of the zero frequency phasor is much greater than those of other phasors. The zero frequency phasor determines the DC offset level due to thermal background in the synthesized image. Not only is this background of little astronomical interest, it is also a major source of Poisson fluctuations. Therefore in this paper we prefer the former method  and consider only the inversion without total counts as in radio astronomy.   We analyze the SNRs of space interferometers in three steps. First, we formulate the SNRs of $^nC_2$ and $^nC_n$ interferometers  in the presence of source shot noise, thermal background shot noise and detector read noise in \\S2 and \\S3 respectively. These formulae cover optical and infrared interferometers. Second, we consider spectrally broad light for which fringe dispersion becomes an issue. Fringes of the $^nC_2$ interferometer can always be dispersed while those of the $^nC_n$ interferometer can be dispersed only in the case of a one dimensional baseline configuration. We discuss the tradeoff between the $^nC_2$ and $^nC_n$ interferometers including the dimensionality of the baseline configuration in \\S4. Third, we apply the results of our analysis to planned interferometers. We present the sensitivities of a TPF-like interferometer and a Darwin-like one in \\S5. Finally in \\S6 we discuss the capability of the TPF-like interferometer from the point of view of deep galaxy count. ",
        "conclusions": "In this paper, we have studied the sensitivities of space borne infrared interferometers for the purpose of general synthesis imaging. Our primary conclusions follow.   (1)  We have derived the expressions for the SNRs of $^nC_2$ and $^nC_n$ interferometers in the presence of source shot noise, thermal background shot noise, and detector read noise.  (2) We have investigated the tradeoff between the $^nC_2$ and $^nC_n$ interferometers for a two dimensional baseline configuration and found that the $^nC_2$ geometry is suitable for a long baseline interferometer, while the  $^nC_n$ geometry is fitted to a short baseline deployable interferometer. For a one dimensional baseline configuration, we have found that the $^nC_n^\\prime$ interferometer is superior especially in the case of read noise limited observations.  (3) We have presented the detection limits of separated spacecraft interferometers like TPF and Darwin in the presence of the natural thermal background as functions of wavelength. The comparison of a TPF-like interferometer with NGST, SPICA and FIRST has revealed that at $\\lambda > 100$ $\\mu$m, an interferometer  is more sensitive than SPICA or FIRST. This is true with only radiation cooling of the apertures.    We thank David Saint-Jacques and the anonymous referee for comments on the manuscript. TN is supported by Grant-in-Aid for Scientific Research of the Japanese Ministry of Education, Culture, Sports, and Science (No. 10640239).  \\appendix"
    },
    "0011/astro-ph0011415_arXiv.txt": {
        "abstract": "We have conducted ultra-deep optical and deep near-infrared observations of a field around the $z=1.226$ radio-quiet quasar 104420.8+055739 from the Clowes-Campusano LQG of 18 quasars at $z\\sim1.3$ in search of associated galaxy clustering. Galaxies at these redshifts are distinguished by their extremely-red colours, with $I-K>3.75$, and we find a factor $\\sim$11 overdensity of such galaxies in a $2.25\\times2.25\\,{\\rm arcmin}^{2}$ field centred on the quasar. In particular, we find 15--18 galaxies with colours consistent with being a population of passively-evolving massive ellipticals at the quasar redshift. They form `fingers' in the $V-K/K$, $I-K/K$ colour-magnitude plots at $V-K\\simeq6.9$, $I-K\\simeq4.3$ comparable to the red sequences observed in other $z\\simeq1.2$ clusters. We find suggestive evidence for substructure among the red sequence galaxies in the $K$ image, in the form of two compact groups, 40 arcsec to the north, and 60 arcsec to the south-east of the quasar. An examination of the wider optical images indicates that this substructure is significant, and that the clustering extends to form a large-scale structure  2--3 $\\mpc$ across. We find evidence for a high ($\\gsim50\\%$) fraction of blue galaxies in this system, in the form of 15--20 `red-outlier' galaxies with $I-K>3.75$ and $V-I<2.00$, which we suggest are dusty, star-forming galaxies at the quasar redshift. Within 30 arcsec of the quasar we find a concentration of blue ($V-I<1$) galaxies in a band that bisects the two groups of red sequence galaxies. This band of blue galaxies is presumed to correspond to a region of enhanced star-formation. We explain this distribution of galaxies as the early-stages of a cluster merger which has triggered both the star-formation and the quasar. ",
        "introduction": "Quasars are sparsely distributed across the universe, with the majority appearing to be unrelated to one another. However quasar clustering has been observed on large scales in the form of Large Quasar Groups (LQGs) (e.g. Webster 1982; Crampton, Cowley \\& Hartwick 1987,1989; Clowes \\& Campusano 1991,1994; Graham, Clowes \\& Campusano 1995), where 4--25 quasars form structures 100--200$\\mpc$ across. They are thus comparable in size to the largest structures seen at the present epoch, such as the `Great Wall' of galaxies, but are seen at an earlier epoch with \\mbox{$0.4\\lsim z\\lsim2$}. It has been suggested \\cite{komberg94} that LQGs represent the progenitors of these local large-scale structures, not only because of their comparable sizes, but also because of a claim that the comoving number densities of LQGs and local superclusters are comparable \\cite{komberg96}.  Observations of the galaxy environments of individual quasars lends credence to the hypothesis that LQGs trace superclusters. Quasars at $z\\sim0.5$ are known to favour young, low to moderately rich clusters, with richness increasing with both redshift and radio-loudness \\cite{ellingson91}. At higher redshifts, a survey of 31 \\mbox{$1\\lsim z\\lsim 2$} radio-loud quasars \\cite{hall98} indicates that they are on average found in moderately rich ($R_{Abell}\\approx1.5$) clusters. Optical and narrow-band (O\\,{\\sc ii}) observations of fields centred on 11 $z\\sim1.1$ quasars from the Crampton \\etal \\shortcite{crampton89} LQG show an excess of galaxies around all but one of the quasars (Hutchings, Crampton \\& Johnson 1995, hereafter HCJ; Hutchings, Crampton \\& Persram 1993, hereafter HCP). In particular they find excesses of blue and emission-line galaxies around several of the quasars indicating regions of significant star-formation, a phenomenon probably related to the fact that both quasars and star-formation require the disruption of the galaxy causing a collapse of significant amounts of gas. They find no evidence for a coherent structure connecting the quasars, but star-formation is likely to be localised, rather than occurring simultaneously across large scales, as a cluster mass halo collapses or during a cluster merger.  To determine whether LQGs trace superclusters requires observations capable of identifying the quiescent galaxies that mark out any associated supercluster, in particular the massive ellipticals which dominate cluster cores. These are both the most luminous and the reddest galaxies in nearby clusters, and form a homogeneous population with very tight linear colour-magnitude (C-M) relations known as {\\em red sequences} \\cite{bower92}, indicative of old (12--13 Gyr) stellar populations. This red sequence has been followed in the optical for clusters out to $z\\sim0.9$ (Arag\\'{o}n-Salamanca \\etal 1993; Stanford, Eisenhardt \\& Dickinson 1998) and is consistent with the passive evolution of galaxies that formed in a monolithic collapse at $z\\gsim3$ \\cite{eggen62}. At $z\\gsim1$ these galaxies should be characterised by extremely-red optical-NIR colours ($I-K\\simeq4$) as the 4000\\AA$\\,$ break is redshifted into the $I$ band. Good contrasts of $z\\gsim1$ clusters over the background are possible as field $I-K\\simeq4$ galaxies appear rare: the Hawaii K-band survey finds no galaxies with $I-K>4$ for $K<18$ over an area of $86.7\\,{\\rm arcmin}^{2}$ (Cowie \\etal 1994; Songaila \\etal 1994).  Several $z\\simeq1.2$ clusters have been found by searching for galaxies with $I-K\\simeq4$ in fields around both targeted high-redshift AGN (Dickinson 1995, hereafter D95; Yamada \\etal 1997; Tanaka \\etal 2000a, hereafter T00a) and regions of extended X-ray emission (Stanford \\etal 1997, hereafter S97; Rosati \\etal 1999, hereafter R99). These galaxies are found to have the optical-NIR colours ($I-K\\simeq4, R-K\\simeq6$) expected for passively-evolving galaxies which are 2--3 Gyr old, in good agreement with the monolithic collapse model predictions. D95 also finds a red sequence for the $z=1.2$ cluster around the radio-galaxy 3C324 with $R-K\\simeq5.9$ and an rms scatter of only 0.07 mag, suggesting that these galaxies formed within 300 Myr of one another.  In this study we aim to establish the galaxy environment of LQGs in a manner unbiased with respect to galaxy type. We have targeted a \\mbox{$30\\times30\\,{\\rm arcmin}^{2}$} field containing three quasars from the $z\\simeq1.3$ Clowes \\& Campusano LQG (Fig.~\\ref{LQG}) for ultra-deep optical observations complemented with deep NIR observations of selected subfields, capable of detecting the passively-evolving galaxies that should mark out any coherent large-scale structures. We present here the first results from this study, using ultra-deep $V$ and $I$ images and a $K$ image centred on the z=1.226 radio-quiet (not detected at the 1mJy level by the FIRST VLA 20cm survey) quasar 104420.8+055739. We have detected clustering of galaxies with the extremely-red colours expected of quiescent ellipticals at the quasar redshift. Throughout the paper we adopt $q_{0}=0.5$ and \\mbox{${\\rm H}_{0}=100h\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$}.  \\begin{figure} \\psfig{figure=figure1.eps,width=4.4in} \\caption{The Clowes-Campusano LQG with each quasar labelled by its redshift. The \\mbox{$30\\times30\\,{\\rm arcmin}^{2}$} region targeted for BTC imaging is indicated by the box, and the quasar for which we have $K$ imaging is circled. The boundaries of the plot match the boundaries of the AQD survey of ESO/SERC field 927 (Clowes \\& Campusano 1991,1994; Clowes, Campusano \\& Graham 1999). Note that this plot and Figures~\\ref{spatial}, \\ref{densitymap} and \\ref{lss} have east to the right.} \\label{LQG} \\end{figure} ",
        "conclusions": "We have found a $3.5\\sigma$ excess of $K<19$ galaxies in the \\mbox{$2.25\\times2.25\\,{\\rm arcmin}^{2}$} field around the z=1.226 radio-quiet quasar 104420.8+055739 from the Clowes-Campusano LQG. The excess is due entirely to a factor $\\sim$11 overdensity of red ($I-K>3.75$) galaxies, which must have $z\\gsim0.8$ to explain their colour. In particular, we find cluster red sequences of 15--18 galaxies, in the $I-K$ and $V-K$ against $K$ C-M diagrams at $I-K\\simeq4.3, V-K\\simeq6.9$, comparable in both colour and magnitudes to red sequences observed for other $z\\simeq1.2$ clusters. These red sequences suggest a population of massive ellipticals at the quasar redshift. In the area of the $K$ image we find tentative evidence of substructure amongst these galaxies, with two apparent groups 40 arcsec to the north and 60 arcsec to the south-east of the quasar. An examination of the optical images over a \\mbox{$7\\times7\\,{\\rm arcmin}^{2}$} area indicates that this substructure is significant at the 2\\% level, and that the clustering extends well beyond the $K$ image, forming a large-scale structure \\mbox{2--3$\\mpc$} across. The overall structure is suggestive of being the early stages of formation of a cluster from the progressive coalescence of subclusters.  Only half of the excess of $I-K>3.75$ galaxies is accounted for by the red sequence members, and we find comparable numbers (15--20) of `red outlier' galaxies  with both $I-K>3.75$ and $V-I<2.00$, which appear to fit neither the passively-evolving nor exponentially-decaying SFR galaxy models (Fig.~\\ref{modelcols}).  Although, all we can say is that they are likely to be dusty star-forming galaxies at $1\\lsim z\\lsim2$, given that such galaxies are found around other $z\\simeq1.2$ clusters (e.g. T00; Kajisawa \\etal 1999; 2000), and given their comparative rarity in field regions, it seems reasonable to assume that they are associated with the cluster. This would suggest that the Butcher-Oemler effect observed in intermediate-redshift clusters increases in strength to higher redshifts, with $\\gsim50\\%$ of likely cluster members exhibiting the blue colours of recent star-formation.  We also find a concentration of blue ($V-I<1$) galaxies within 30 arcsec $(130\\kpc)$ of the quasar, with many having the red $I-K$ colours of $z\\gsim1$ galaxies, which suggests that the quasar lies in a region of enhanced star-formation, in agreement with the results of HCJ. This concentration appears extended in such a way as to separate the two groups of red galaxies, and it is also notable how the blue galaxies appear to avoid the centres of red galaxy clustering. Some foreground contamination is likely, but given that such concentrations appear common around LQG quasars at this redshift, and the highly significant spatial interrelation between the blue and red galaxies, then we can be reasonably confident that this concentration is real, and is associated with the quasar. However spectroscopic observations of these and the other cluster candidates will be required to confirm cluster membership, and to provide more quantitative information about their star-formation histories, such as their approximate ages and current star-formation rates.  \\subsubsection{A Possible Cluster Merging Event}  The relative distribution of red and blue galaxies can be explained if what we are witnessing is the early-stages of merger of the two clusters of red galaxies, which has triggered both the band of enhanced star-formation and the quasar itself. A comparable distribution has been observed for the Coma cluster \\cite{caldwell93} with a band of post-starburst or `E+A' galaxies located between the Coma cluster centre and a secondary X-ray peak. Dynamical studies \\cite{burns94} indicated that these galaxies had passed through the centre of the Coma cluster about 2 Gyr ago, coincident with the epoch of starbursting predicted from the spectra of the post-starburst galaxies. Unusually high blue galaxy fractions have been observed for a number of low-redshift clusters with bimodal X-ray surface brightness profiles \\cite{metevier2000}, implying that cluster mergers can induce starbursts simultaneously in a large fraction of cluster galaxies, and they could be a major contributor to the Butcher-Oemler effect. Several mechanisms have been suggested that could cause firstly the triggering and then the termination of a secondary burst of star-formation in a galaxy, as a subcluster passes through a cluster. These include ram pressure from the ICM \\cite{evrard91}, shocks due to collisions between the two ICMs \\cite{roettiger96}, and the effect of close galaxy encounters \\cite{moore96}. As the `band' of blue galaxies is likely to be undergoing or has recently undergone star-formation, and because the two groups are relatively close together, we suggest that this system is being observed at an earlier epoch of the cluster merger process than Caldwell \\etal (1993), either just before or just after core passage, and that the star-formation has been triggered by the interaction of galaxies with the shock fronts produced by the collision of the two ICMs (see Roettiger \\etal 1996).  Cluster merging events are predicted to be relatively common at high redshifts ($z\\gsim1$) in hierarchical clustering models (e.g. Press \\& Schechter 1974; Bahcall, Fan \\& Cen 1997; Percival \\& Miller 1999). Examples of possible merging clusters at high redshifts are the CL0023+00423 groups at z=0.8274 and 0.8452 which, according to a dynamical study, have a 20\\% chance of merging \\cite{lubin2000}, and the R99 and S97 clusters which are separated by only $2.5\\mpc$.  \\subsubsection{Comparison with Other Work}  In a study of 7 radio-loud quasars at $1.0<z<1.6$, S\\'{a}nchez \\& Gonz\\'{a}lez-Serrano (1999) find excesses of faint ($B>22.5$ and $R>22.0$) galaxies on scales of $r<170$ arcsec and $r<35$ arcsec around the quasars, whose numbers, magnitudes and angular extensions are compatible with being clusters of galaxies at the quasar redshifts. In particular however, they find that the quasars are in general not located at the peaks of the density distribution, but are some 40--100 arcsec from them, located on the cluster peripheries. This result, and our own, is understandable in the framework of quasar activity being triggered by the infall of gas onto a seed black hole. Firstly, galaxies in the centres of clusters have previously lost most or all of their gas, by having had it stripped off by ram-pressure from the intra-cluster medium or by tidal forces from close encounters with other galaxies, or by consuming the gas in a starburst during its first infall into the cluster. Secondly, the encounter velocities for galaxies in the cluster cores greatly exceed the internal velocity distribution of the galaxies, making galaxy mergers much less effective at triggering nuclear activity \\cite{aarseth80}.  In a comparable study of the galaxy environment of the radio-loud quasar 1335.8+2834 from the $z\\sim1.1$ Crampton \\etal (1989) LQG, Tanaka \\etal (2000a) obtain results which have several similarities to our own. This quasar had been part of the HCJ study, and was known to lie in a band of blue and emission-line galaxies. Using deep $R,I$ and $K$ observations, a number of extremely-red objects with the colours of passively-evolving galaxies at the quasar redshift were found, forming a cluster which lies to one side of both the quasar and the band of blue and emission-line galaxies. They also find a similar population of `red outliers' and estimate the blue galaxy fraction as 60--80\\%. They also find an indication that this cluster is part of a larger structure with groupings of optically-red (\\mbox{$R-I>1.3$}) galaxies, similar to the clustering near the quasar, found across the \\mbox{$8\\times8\\,{\\rm arcmin}^{2}$} (\\mbox{$4.7\\times4.7\\,h^{-2}\\,{\\rm Mpc}^{2}$} at z=1.086) $R,I$ optical images.  In a wide-field ($48\\times9\\,{\\rm arcmin}^{2}$) optical imaging survey of the 1338+27 field containing 5 quasars from the Crampton \\etal (1989) LQG, Tanaka \\etal (2000b) detect significant clustering of faint red galaxies with $I>21$ and \\mbox{$R-I>1.2$}. These galaxies are concentrated in 4--5 clusters forming a linear structure of extent $\\sim10\\mpc$ that is traced by the group of quasars, although only the one radio-loud quasar of T00a appears to be directly associated with any of the rich clusters. The immediate environments of the other four LQG quasars, all radio-quiet, appear relatively poor in terms of red galaxies, although three have excesses of blue and emission-line galaxies (HCJ95) indicating that they are located within regions of enhanced star-formation. This is the clearest evidence yet that LQGs trace large-scale structure, even if the majority of the member quasars are only directly associated with regions of enhanced star-formation, rather than rich clusters.  \\subsubsection{Interpretation - Mechanisms for Quasar Formation}  As several of the HCJ quasars were found in regions of enhanced star-formation, and as both the quasar of T00a and this paper are located in `bands' of enhanced star-formation, in between, or on the peripheries of, clusters, we propose a causal link between the quasar and star-formation whereby both are triggered by the same mechanism: the interaction between the galaxy and the intra-cluster medium. If a galaxy can be disrupted sufficiently by its passage through the ICM to cause it to undergo starbursting, then if it also contains a supermassive black hole, enough gas may be channelled onto the nucleus to trigger a phase of quasar activity. Such a mechanism explains the finding of quasars in regions of enhanced star-formation more naturally than the galaxy merger model, as it allows many galaxies to be affected simultaneously, although there is good evidence that a large fraction of quasars have been triggered by galaxy mergers. If there is a connection between the quasar and star-formation activation, then it is likely that these quasars are found preferentially in clusters with high blue galaxy fractions, and that these are not representative of $z\\simeq1.2$ clusters as a whole.  This work and previous studies show that searching for sources with the $I-K>3.75$ colours characteristic of quiescent galaxies at $z\\gsim1$ is an efficient means of locating $z\\gsim1$ clusters. By adding a second optical band it is then possible to obtain qualitative information on the star-formation history of these galaxies. Using this information along with the relative spatial distribution of quiescent and star-forming galaxies, a more complex picture arises in which the evolution of galaxies, quasars and clusters are all interrelated."
    },
    "0011/astro-ph0011286_arXiv.txt": {
        "abstract": "A detailed photometric and spectroscopic analysis of cool ($\\Te\\ \\lta 12,000$~K) white dwarf stars is presented. The sample has been drawn from the Yale Parallax Catalog and from a proper-motion survey in the southern hemisphere. Optical $BVRI$ and infrared $JHK$ photometry, as well as spectroscopy at \\halpha, have been secured for a sample of 152 white dwarfs. The discovery of seven new DA white dwarfs, two new DQ white dwarfs, one new magnetic white dwarf, and three weak magnetic white dwarf candidates, is reported. Our sample also identifies 19 known or suspected double degenerates.  The photometric energy distributions, the \\halpha\\ line profiles, and the trigonometric parallax measurements are combined and compared against the predictions of model atmosphere calculations to determine the effective temperature and the radius of each object in the sample, and also to constrain the atmospheric composition. New evolutionary sequences with carbon/oxygen cores with thin and thick hydrogen layers are used to derive stellar masses and ages. The results are used to improve our understanding of the chemical evolution of cool white dwarfs.  We confirm the existence of a range in effective temperature between $\\sim$5000 and 6000~K where almost all white dwarfs have hydrogen-rich atmospheres. Our sample shows little evidence for mixed H/He white dwarfs, with the exception of two helium-rich DA stars, and four (possibly five) C$_2$H white dwarfs which have been interpreted as having mixed H/He/C atmospheres. The observed sequence of DQ stars is found to terminate abruptly near 6500~K, below which they are believed to turn into C$_2$H stars. True DC stars slightly above this temperature are found to exhibit hydrogen-like energy distributions despite the lack of \\halpha\\ absorption features. The mean mass of our complete sample is 0.65 \\msun\\ with a dispersion of $\\sigma\\sim0.20$ \\msun.  Attempts to interpret the chemical evolution of cool white dwarfs show the problem to be complex. Convective mixing is called upon to account for the increase of the non-DA to DA ratio below 12,000~K, as well as the reappearance of helium-rich stars below $\\sim 5000$~K. The possible presence of helium in cool DA stars, the existence of the non-DA gap, and the nature of the peculiar DC stars are also explained in terms of convective mixing, although our understanding of how this mechanism works needs to be revised in order to account for these observations. Given this chemical evolution uncertainty, it is not clear whether thick or thin hydrogen layer models should be used to determine cooling ages. The oldest object in our sample is $\\sim$7.9~Gyr or $\\sim$9.7~Gyr old depending on whether thin or thick hydrogen layer models are used, respectively. ",
        "introduction": "Cool white dwarfs have recently gained attention due to the discovery of faint ($V \\sim 26-28.5$) objects in the Hubble Deep Field Survey \\citep[see][and references therein]{hansen98} which have been interpreted by \\citet{hansen98, hansen99} as being cool ($\\Te\\ \\lta 3000$~K) and old ($t\\ \\gta 11$ Gyr) hydrogen atmosphere white dwarfs. Hansen has shown that the location of these objects {\\it blueward} of the white dwarf cooling sequence in the \\mv\\ vs $(V-I)$ color-magnitude diagram is the result of the extremely strong \\htwo\\ molecular absorption features which extend well into the optical regions of their energy distributions. Hence, even though hydrogen-rich white dwarfs become redder with decreasing effective temperature down to $\\Te\\sim 4000$~K \\citep[see, e.g., Fig.~6 of][]{bwb95}, even cooler stars actually become bluer. Note that the most recent photometric and spectroscopic analysis of cool white dwarfs in the local Galactic disk \\citep[hereafter BRL]{brl} has shown no evidence of white dwarfs cooler than 4000~K.  In addition, \\citet{hansen99} has also discussed the importance of using detailed stellar atmospheres as boundary conditions for calculating cooling white dwarf models. For instance, the cooling age of a 0.6 \\msun, $\\Te\\sim 5000$~K white dwarf can differ by $\\sim 1.5$ Gyr depending on whether the object has a hydrogen- or a helium-dominated atmosphere. Given this sensitivity, it is important to determine precisely the atmospheric composition of cool white dwarfs before any attempt to determine their age is made.  To complicate things even more, white dwarfs are not expected to preserve their hydrogen or helium atmospheric signature throughout their evolution. Indeed, there is both theoretical and observational evidence which indicates that the chemical composition of white dwarfs is likely to change many times during their lifetime due to competing mechanisms (e.g., gravitational settling, convective mixing, convective dredge-up, accretion from the interstellar medium, radiative acceleration) which may affect the composition of the outer layers of these stars as they cool  \\citep[see][for a review on the subject]{fw87}. This is even more important for cool white dwarfs where convective mixing of a thin ($M_{\\rm H}/M_\\star\\ \\lta10^{-6}$) hydrogen atmosphere with the more massive convective helium envelope can turn a hydrogen-rich atmosphere into a helium-dominated atmosphere anywhere between 12,000~K and 5000~K. The exact temperature at which convective mixing occurs --- or whether it will occur at all --- depends on the thickness of the hydrogen layer (the thicker the hydrogen layer, the cooler the mixing temperature). Furthermore, BRL have clearly demonstrated that the chemical evolution of cool white dwarfs is even more complex than previously anticipated, and that several mechanisms --- some of them still unknown --- are at play to alter the chemical constituent of their atmospheres as they cool. Consequently, the age of any cool white dwarf is not only strongly dependent on its atmospheric composition, as pointed out by \\citet{hansen99}, {\\it but also on its chemical evolution history}. For instance, a recently mixed helium-rich atmosphere white dwarf would have an integrated cooling age of its immediate predecessor --- i.e.~that of a hydrogen-rich atmosphere white dwarf. Stellar ages and chemical evolution can thus not be easily disentangled. This unfortunate consequence stresses the importance of better understanding the chemical evolution of cool white dwarf stars.  For several years now, we have been pursuing our comprehensive investigation of cool white dwarf evolution. We presented in BRL photometric and spectroscopic observations of 110 cool white dwarfs, and analyzed them with model atmosphere calculations appropriate for cool white dwarfs with pure hydrogen and pure helium compositions, as well as with mixed H/He compositions. Our fitting technique consists of comparing the observed energy distributions in the optical $BVRI$ and infrared $JHK$ with those predicted from model atmospheres in order to determine the effective temperature of the star. When available, trigonometric parallax measurements are used to derive the radius of the star, and hence the mass through the mass-radius relation for white dwarfs. Finally, high signal-to-noise spectroscopy at \\halpha\\ is used to determine, or at least constrain, the chemical composition of their atmospheres. The main results of the BRL analysis are summarized here together with our current understanding of the physical models which have been proposed to account for the observations.  Most cool white dwarfs have energy distributions which are consistent with either pure hydrogen or pure helium atmospheric compositions, with little evidence for objects with mixed H/He compositions. In general, the energy distributions of DA (hydrogen-line) stars are consistent with pure hydrogen atmospheres, while those of non-DA stars are consistent with pure helium atmospheres. Exceptions include a few extremely cool non-DA stars with hydrogen-dominated atmospheres which are simply too cool to exhibit any hydrogen absorption features, the so-called C$_2$H stars whose peculiar energy distributions and strong molecular absorption bands have been interpreted as the result of mixed H/He/C compositions with $\\nh \\sim 10^{-2}$ \\citep{schmidt95}, and a few other objects which are most likely unresolved binaries (see BRL and references therein). Note that hydrogen lines are observed in DA stars as cool as $\\Te\\sim 4500$~K, that is, almost at the end of the observable white dwarf cooling sequence of the local Galactic disk.  BRL have also shown that the temperature distribution of DA and non-DA stars is not homogeneous. In particular, even though non-DA stars are abundant above $\\Te\\sim 6000$~K, most white dwarfs below this temperature, down to $\\Te\\sim 5000$~K, are of the DA type (the so-called ``non-DA gap''). Moreover, BRL have identified a group of objects which lie close to, and above, $\\Te\\sim 6000$~K whose energy distributions are better reproduced with pure hydrogen models, even though their optical spectra are completely featureless; hydrogen atmosphere white dwarfs in this particular range of temperature should exhibit strong hydrogen lines. BRL proposed an exotic model to account for both the strange nature of these objects and the existence of the non-DA gap.  A warmer, pure helium atmosphere white dwarf slowly accretes hydrogen from the interstellar medium as it cools. Most atomic levels of the hydrogen accreted are quenched as a result of the high atmospheric pressures which characterize cool, helium-rich atmospheres. In particular, the bound level of \\hm\\ is destroyed, and the corresponding bound-free \\hm\\ opacity --- the most important opacity source in hydrogen-rich atmospheres --- is significantly reduced. Consequently, the accretion of hydrogen onto a helium-dominated atmosphere would have little effect on the opacity and thus on the atmospheric pressure. BRL proposed that large amounts of hydrogen could be {\\it hidden} through this simple accretion/quenching mechanism. They also suggested that at some point, the free-free \\hm\\ opacity would become large enough for the atmospheric pressure to drop below the threshold value at which the \\hm\\ ion starts to recombine. The star then rapidly turns into a normal hydrogen-rich white dwarf. BRL suggested that the group of peculiar white dwarfs discussed above could represent such objects with a mixture of helium and quenched hydrogen, in the process of turning into normal DA stars, and that the non-DA gap could be accounted for by the sudden transformation of apparently helium-rich atmosphere white dwarfs into hydrogen-dominated atmospheres.  More recently, \\citet{malo99} have explored more quantitatively the model proposed by BRL. The results of \\citeauthor{malo99} indicate that the accretion of even minute amounts of hydrogen onto a helium-dominated atmosphere, $\\log \\nh\\ \\gta 10^{-6}$, provide enough free electrons --- through the thermal ionization of hydrogen --- to increase the He$^-$ opacity significantly, which in turn, reduces the atmospheric pressure below which quenching effects on hydrogen become negligible. The presence of hydrogen therefore has an {\\it indirect} effect on the total opacity through the free electrons it provides to the plasma. On the basis of their analysis, \\citeauthor{malo99} rejected the model proposed by BRL. Given these results, the nature of the peculiar objects identified by BRL and the existence of the non-DA gap remain unexplained.  \\citet{hansen99} provided an alternative explanation for the non-DA gap. The gap is evident in his Figure 13 even though he restricted his sample of objects to that of \\citet{ldm88} rather than using the more complete sample studied by BRL, where the gap is even more obvious. Hansen has shown that the concentration of non-DA stars just below $\\Te\\sim 5000$~K have nearly the same age as the concentration of DA stars around 6000~K. On the basis of this conclusion alone, Hansen suggested that the non-DA gap alluded to by BRL is only a result of the difference of cooling timescales between hydrogen-dominated and helium-dominated atmospheres. However, this interpretation cannot be correct since it neglects the fact that non-DA stars are also observed in large numbers {\\it above} the gap at $\\Te\\ge 6000$~K.  In other words, non-DA stars are seen both above and below the gap, and the high-temperature end of the gap does not reflect a simple termination of the cooling sequence of degenerates with helium-dominated atmospheres.  Unless helium-rich white dwarfs evolve more rapidly {\\it through} the gap than above or below the gap, which is not the case, an explanation based solely on differences in cooling rates does not hold. Obviously, a chemical evolution from one spectral type to another must be invoked.  Other results from the BRL analysis support the idea that some physical mechanism, yet unidentified, is at play near the blue edge of the non-DA gap. For instance, the most widely accepted explanation for the presence of carbon in the photospheric layers of DQ stars is that carbon is convectively dredged-up from the core, as demonstrated quantitatively by \\citet{pelletier86}. These calculations show that the extent of the carbon diffusion tail is reduced at low effective temperature due to the partial recombination of carbon, and smaller carbon abundances are predicted in DQ stars with lower effective temperature, in apparent agreement with the observations \\citep[see Fig.~4 of][]{pelletier86}. As pointed out by BRL, however, the coolest DQ stars have the strongest observed C$_2$ molecular bands, and cooler objects with even smaller abundances are well within the observational limits and should be easily detectable. The BRL analysis has revealed instead the existence of an abrupt cutoff in effective temperature at $\\Te\\sim 6500$~K below which DQ stars are no longer found. The presence of carbon in stars below this temperature threshold is in the form of what has been interpreted as C$_2$H molecular bands by \\citet{schmidt95}. The most obvious explanation proposed by BRL is that DQ stars turn into C$_2$H stars below $\\Te\\sim 6500$~K. The mechanism proposed by BRL by which the hydrogen required to form the C$_2$H molecule suddenly appears in these objects relies on their composition transition model, which was later discarded by \\citet{malo99}. Hence, even though the DQ to C$_2$H transition remains the most likely explanation for the existence of the abrupt cutoff of the DQ star phenomenon, the physical mechanism by which this transition occurs needs further investigation.  BRL have also shown that below the non-DA gap near $\\Te\\sim 5000$~K, non-DA stars reappear in large number. The most likely explanation for this phenomenon is the convective mixing of the relatively thin surface hydrogen layer with the deeper and more massive helium envelope. Detailed envelope calculations indicate that the bottom of the hydrogen convection layer reaches its deepest point at a certain temperature \\citep[see, e.g.,][]{tassoul90}; white dwarfs with hydrogen layers thicker than this maximum depth point can never mix and will remain hydrogen-rich forever. The canonical value of the lowest possible mixing temperature is $T_{\\rm mix}\\sim6000$~K, which corresponds to a hydrogen layer mass of $M_{\\rm H}/M_\\star\\sim10^{-6}$. This is $\\sim 1000$~K hotter than the temperature required to account for the BRL observations. New envelope models discussed by BRL which include realistic model atmospheres as boundary conditions reduce the mixing temperature limit well down to 5000~K, however. Hence, convective mixing may well represent the correct interpretation for the reappearance of non-DA stars below the non-DA gap.  The continued presence of DA stars requires that at least some stars have hydrogen layers thicker than $M_{\\rm H}/M_\\star\\sim10^{-6}$.  Following the BRL analysis, \\citet{lrb98} have gathered new photometric and spectroscopic data for the 43 cool, low-luminosity, high--proper-motion white dwarfs in the \\citet{ldm88} sample.  This sample was analyzed with the techniques developed by BRL in order to improve the determination of the cool white dwarf luminosity function and the estimate of the age of the local Galactic disk. BRL discuss at length the shortcomings of the previous estimates by \\citet{ldm88}, and in particular the lack of precise atmospheric parameter determinations as well as accurate bolometric corrections for the stars in the sample. The improved luminosity function of \\citet{lrb98} yields an age of the local region of the Galactic disk of $8\\pm1.5$ Gyr, with the uncertainty being mainly due to the core composition, the effect of separation on crystallization, and the unknown mass of the outer helium envelope.  The \\citet{ldm88} sample is being extended to a lower proper-motion limit and a new luminosity function is being determined which supports the earlier work \\citep{ldhl99}.  As mentioned in BRL, these earlier efforts represent only the first step towards a better understanding of the chemical evolution of cool white dwarfs. Even though the BRL analysis has provided us with a wealth of new observational and theoretical results, many questions have remained unanswered, and the problems to be solved have multiplied. The success of our analysis has prompted us to extend our study to a larger and more complete sample of cool white dwarfs. We thus present in this paper a detailed analysis of all cool ($\\Te\\ \\lta 12,000$~K) white dwarfs with trigonometric parallax measurements available to us. Such measurements are of particular importance to derive stellar radii, and hence stellar masses. This new sample is comprised of 152 white dwarfs for which masses can be derived. In contrast, only 67 out of the 110 white dwarfs analyzed by BRL had trigonometric parallax measurements. Furthermore, BRL concentrated their effort towards cool ($\\Te\\ \\lta 7000$~K), non-DA white dwarfs, while our new sample is distributed over a wider range of effective temperature and spectral types.  In \\S\\ 2, we discuss the improvements in the model atmospheres and cooling models since our most recent analysis. In \\S\\ 3, we present our new photometric and spectroscopic observations, which, combined with our earlier observations, are analyzed in \\S\\ 4. The global properties of our sample are examined in \\S\\ 5. Our conclusions follow in \\S\\ 6. ",
        "conclusions": "A detailed photometric and spectroscopic analysis of 152 cool white dwarfs with trigonometric parallax measurements was presented. Optical and infrared photometry were combined and compared with an improved grid of model atmosphere calculations to determine the effective temperatures and the solid angles for all objects, which combined with the measured parallaxes yielded the stellar radii. The latter were then converted into stellar masses through the mass-radius relations obtained from new evolutionary models which best characterize cool white dwarfs. These have carbon/oxygen cores, with both thin $M_{\\rm H}/M_\\star = q({\\rm H})=10^{-10}$ and thick $q({\\rm H})=10^{-4}$ hydrogen layers. This parallax sample more than doubles the number of cool white dwarfs with known masses compared to the sample presented in BRL. Furthermore, this new sample is more representative of the population of cool white dwarf stars as it contains, presumably, no bias in favor of any particular spectral type, as opposed to the BRL sample which favored non-DA stars.  One of the main results of our analysis is the confirmation of the existence of a range in effective temperature between roughly 5000 and 6000~K in which almost all white dwarfs have hydrogen-rich atmospheres. The only non-DA stars in this range of effective temperature are the so-called C$_2$H stars which are believed to represent the cooler stage of the DQ sequence, found only above the gap. Clearly, the existence of this gap holds the key to our understanding of the chemical evolution of cool white dwarfs. Unfortunately, our understanding of the physical mechanisms that could alter the atmospheric composition of cool white dwarfs as they evolve are at best sketchy. After the ``phase transition'' model proposed in BRL was convincingly discarded by \\citet{malo99}, very few alternatives exist.  It has become clear from our analysis, however, that convection plays an important role, but that our understanding of what to expect from the result of convective mixing needs to be revised. It is generally believed that when the bottom of the thin hydrogen convection zone deepens and reaches the underlying and more massive convective helium envelope, convective mixing will occur, and that the resulting atmospheric composition will simply be governed by the mass ratio of the hydrogen and helium convection zones. Such arguments lead to predicted abundances after convective mixing of $\\nh=10^{-4}$, $10^{-3}$, $10^{-2}$, and $10^{-1}$, for mixing temperatures of $T_{\\rm mix}\\sim11,000$, 10,000, 8500, and 6500~K, respectively, the mixing temperature depending on the exact thickness of the hydrogen layer. Our analysis has demonstrated that such objects do not exist, at least in our sample. Even the objects which possess mixed hydrogen and helium atmospheres, and in particular Ross 640 and L745$-$46A, have abundances that do not fit this simple convective mixing picture. On the other hand, the definite increase of the ratio of non-DA to DA stars in the temperature range considered in this analysis with respect to that of hotter white dwarfs cannot easily be explained in terms other than convective mixing. Also, convective mixing has been called upon to explain the reappearance of non-DA stars below the non-DA gap near 5000~K.  A different picture is now starting to emerge from our findings, and those of other analyses, in which convective mixing {\\it processes} may lead to predicted chemical abundances that differ from the standard model. First, our comparison of atmospheric parameters determined from photometry and spectroscopy reveals that a mass discrepancy exists in several DA stars, and that this discrepancy can be resolved if we allow some amount of helium to be present in their atmospheres. The required abundances close to $\\nhe=1$ are clearly at odds with the standard convective mixing model predictions, as noted also by \\citet{bergeron90}. Note that helium is spectroscopically invisible in the temperature regime considered here and cannot be detected directly. It is possible to conceive that helium is being brought to the surface not by convective mixing, but perhaps by convective overshooting of the helium convection zone into the hydrogen layer. The bottom of the hydrogen convection zone would then carry some of this helium to the surface, in a very similar fashion to the way carbon is being brought to the photosphere of DQ stars, with the exception that ordinary diffusion is responsible in this case for polluting the bottom of the helium layer with carbon.  A variance of the convective model that could account for the existence of the non-DA gap has also been proposed. The question arises when one considers a helium-rich atmosphere on top of which indefinite amounts of hydrogen are accumulated, presumably from interstellar accretion. It is clear that the hydrogen abundance ratio will gradually increase until some point where the homogeneously mixed hydrogen and helium convection zone can no longer be sustained. A stable chemical separation between the hydrogen and the helium layers could then occur, leading to the predominance of DA stars below $\\Te\\sim6000$~K. The implication of this proposed model is that cool white dwarfs must accrete reasonable amounts of hydrogen from the interstellar medium for this mechanism to work. Also, it was proposed that the chemical separation could occur in a non-homogeneous fashion, leading to chemically inhomogeneous atmospheres, similar to the hotter DAB stars \\citep[see][]{koester94, beauchamp93}. Such inhomogeneous atmospheres could perhaps account for the peculiar objects identified in our sample whose energy distributions share both characteristics of hydrogen and helium atmospheres.  Given the complexity of the chemical evolution of cool white dwarf atmospheres revealed in our analysis, it is not clear what evolutionary models should be used to derive cooling ages. Thin and thick hydrogen layer models yield cooling ages which differ by 2 Gyr in some cases. Non-DA stars are probably better characterized by thin hydrogen models; even if they were once DA stars, they probably had thin hydrogen layers anyway in order to eventually mix. But cool DA stars may have thin or thick hydrogen layers. Thick hydrogen layer white dwarfs remain DA forever, but the fraction of such stars is still unknown. The problem becomes even more complex when one attempts to determine the age of the Galactic disk using the luminosity function of cool white dwarfs. The theoretical luminosity functions should in fact reflect the actual population of cool white dwarfs, and should be a combination of thin and thick hydrogen layer models. The problems are greatly simplified for hotter DA, DB, or DO stars. From the results presented here, the oldest white dwarfs in our sample have a minimum stellar age of 8.5 Gyr, but are more likely characterized by models with thick hydrogen layers which yield an age around 9.5 Gyr.  We feel at this point that we have pushed both the observational and theoretical aspects of this cool white dwarf study to their limits, and that more work now needs to be done on the theoretical front to explain the various features discovered in this analysis. Several avenues of investigation have been proposed, most of which are related to our poor understanding of convective mixing."
    },
    "0011/astro-ph0011553_arXiv.txt": {
        "abstract": "s{A massive black hole is present at the centre of our galaxy and inevitably accretes dark matter particles, creating a region of very high particle density. The annihilation rate  is enhanced  with a large number of $e^+e^-$ pairs produced either directly or by sucessive decays of mesons. We evaluate the synchrotron emission (and self-absorption) associated with the propagation of these particles through the galactic magnetic field and give constraints on the values of mass and cross section of the  dark matter particles.} ",
        "introduction": "There is  convincing evidence for the existence of an unseen non-baryonic component in the energy density of the universe. The most promising dark matter candidate appear to be weakly interacting massive particles (WIMPs) and in particular the so-called neutralinos, arising in supersymmetric scenarios  (for a review of particle candidates for dark matter see e.g. \\cite{ellis}). The annihilation of these X-particles would produce quarks, leptons, gauge and Higgs bosons and gluons. In particular $e^+e^-$ pairs are produced either directly or by successive decays of mesons, and lose their energy through synchrotron radiation as they propagate in the galactic magnetic field. This radiation is expected to be greatly enhanced in the proximity of the galactic centre, where the existence of a massive black hole creates a region of very high dark matter particle density and consequently a substantial increase in the annihilation rate and in the ensuing synchrotron radiation.   \\begin{figure}[t] \\epsfxsize=27pc \\epsfbox{absnu.eps} % \\caption{Left panel: $A_{\\nu}$ as a function of frequency for $m_X=1$TeV. The two upper curves correspond to the cross section $\\sigma v=10^{-28}/m_X^2(\\rm GeV)\\,\\rm cm^3s^{-1}$, close to the unitarity limit; the two lower curves correspond to $\\sigma v=10^{-38}/m_X^2(\\rm GeV)\\,\\rm cm^3s^{-1},$ a cross section more typical for wimps. Results for two values of the density profile are shown in each case: $\\gamma=1$, for solid curves and $\\gamma=1.5$ for dashed ones. Right panel: $A_{\\nu}$ as a function of the particle mass for $\\nu$=408MHz, $\\sigma v=10^{-38}/m_X^2(\\rm GeV)\\,\\rm cm^3s^{-1}$ and two values of $\\gamma$. \\label{freq}} \\end{figure} ",
        "conclusions": "We have  shown that present data on the emission from Sgr A*  are compatible with  a wide set of dark matter parameters. The evaluation of synchrotron self-absorption has enabled us to reach an alternative conclusion to an earlier claim of incompatibility of cuspy halos with the existence of annihilating wimp dark matter. We find that the experimental data on the Sgr A* spectrum at radio wavelengths could be explained by synchrotron emission of electrons produced in the annihilation of rather massive dark matter particles, extending up to and beyond TeV masses. This is relevant to a recent study of coannihilations \\cite{boehm}, which  suggests that WIMPs with  $\\Omega_X h^2=0.2$ can extend up to several TeV, as well as to very massive  particles (wimpzillas) that are produced non-thermally in the primordial universe. We also find that the synchrotron emission tends to give somewhat sharper constraints on masses and cross sections than the observed  gamma-ray fluxes and neutrino limits. This situation could change  with more sensitive  gamma-ray and neutrino observations anticipated from  forthcoming experiments."
    },
    "0011/astro-ph0011547_arXiv.txt": {
        "abstract": "We report the discovery of a 40~arcsec long X-ray filament in the core of the cluster of galaxies A\\,1795. The feature coincides with an H$\\alpha+$NII filament found by Cowie et al in the early 1980s and resolved into at least 2 U-band filaments by McNamara et al in the mid 1990s. The (emission-weighted) temperature of the X-ray emitting gas along the filament is $2.5-3\\keV$, as revealed by X-ray colour ratios. The deprojected temperature will be less. A detailed temperature map of the core of the cluster presented. The cD galaxy at the head of the filament is probably moving through or oscillating in the cluster core. The radiative cooling time of the X-ray emitting gas in the filament is about $3\\times 10^8\\yr$ which is similar to the age of the filament obtained from its length and velocity. This suggests that the filament is produced by cooling of the gas from the intracluster medium. The filament, much of which is well separated from the body of the cD galaxy and its radio source, is potentially of great importance in helping to understand the energy and ionization source of the optical nebulosity common in cooling flows. ",
        "introduction": "The rich cluster Abell\\,1795, at redshift $z=0.063$, has been well-studied at optical, radio and X-ray wavelengths. It shows a relaxed X-ray structure (Buote \\& Tsai 1996) with a strong central peak of cooler gas (Fabian et al 1994; Briel \\& Henry 1996; Allen \\& Fabian 1997; Ikebe et al 1999; Allen et al 1999). This is good evidence for a cooling flow (Fabian 1994) which is supported by the presence of strong emission-line nebulosity around the central cD galaxy (Cowie et al 1983; van Breugel, Heckman \\& Miley 1984; Hu, Cowie \\& Wang 1985; Heckman et al 1989) accompanied by excess blue light (Johnstone, Fabian \\& Nulsen 1987; Smith et al 1997; Cardiel, Gorgas \\& Aragon-Salamanca 1998). The blue light is probably due to massive stars (McNamara et al 1996a), with some from young globular clusters (Holtzman et al 1996). Molecular gas in the cD has been detected through molecular hydrogen emission (Falcke et al 1998). The small central radio source 4C\\,26.42 shows high Faraday rotation measures (Ge \\& Owen 1993) which is also a characteristic of a cooling flow (Taylor et al 1991, 1999). The radio source may have triggered some of the star formation in the cD (McNamara et al 1996b).  Hill et al (1988) reported that the cD in A1795 is not at rest in the gravitational potential of the cluster but has a peculiar radial velocity of $365\\kmps$. This was supported by the outer emission-line nebulosity having a blueshift relative to the cD galaxy. Oegerle \\& Hill (1994) later revised the velocity, and the significance of the effect, down to about $150\\kmps$, and argue that a subcluster has merged with the cluster.  Here, we present for the first time Chandra X-ray images of the core of A\\,1795. The high spatial resolution of Chandra allows us to see an 80~kpc long filament of X-ray emission extending SSE from the central galaxy. This filament coincides with a an H$\\alpha$ emission feature discovered by Cowie et al (1983). We discuss possible causes for such a striking feature. ",
        "conclusions": "The Chandra images confirm the short cooling time of the gas in the centre of A1795 and also that the gas is cooler toward the centre. The shortest cooling time is about $3\\times 10^8\\yr$ close to the cD galaxy and along the brightest part of the filament. These results are consistent with the presence of a strong cooling flow in the cluster core.  We have discovered a filament of X-ray emission coincident with the H$\\alpha$ filament of Cowie et al (1983). The X-ray luminosity of the filament, assuming a width of 5 arcsec and length of 40 arcsec (corresponding to a projected length of about 80~kpc, assuming $H_0=50\\kmpspMpc$), is about $2\\times 10^{42}$\\ergps. This is only about an order of magnitude greater than the H$\\alpha+$NII luminosity of the filament. This last luminosity probably converts to a bolometric line luminosity comparable to (or possibly greater than) the X-ray luminosity when all optical and UV emission lines are taken into account. If there is no heating balancing cooling then the mass cooling rate in the filament alone is about $10\\Msunpyr$. The X-ray brightest region in the cD galaxy also coincides with the H$\\alpha+$NII brightest region there.  The age of the filament given its projected length and radial velocity of $\\sim 150$\\kmps corresponds to $5\\times 10^8$\\yr, if it is at an angle of 45 deg to the line of sight. The timescale is of course less if the filament lies closer to the plane of the Sky. This raises the possibility that its length is constrained by cooling. We have found (Fig.~4) that the coolest part of the X-ray filament is near the end opposite to the cD galaxy. The straightness of the filament argues against the intracluster medium being very turbulent.  We note that the radial velocity of the cD relative to the velocity dispersion of the cluster ($\\sigma = 920\\kmps$, Oegerle \\& Hill 1994) is small. This suggests that if the angle of its motion to the line of sight is low then the cD must be very close to the turnaround point in its orbit. An approximate estimate for the difference between the position of the cD, $\\Delta r$, from the turnaround radius of its (radial) orbit, $r$, assuming that the core is of constant density out to radius $a,$ is given by $\\Delta r/r\\approx 0.33 (v/\\sigma)^2(a/r)^2.$ If the angle is large (close to the plane of the Sky) then the restriction is much less and the true velocity much higher.  Although the filament is clear in both the X-ray and H$\\alpha+$NII images the brightness variations along it do not compare well. The brightest part of the H$\\alpha+$NII filament about 25 arcsec S of the centre of the cD galaxy corresponds to a minimum in the X-ray surface brightness. The reverse is true of the X-ray bright part almost 20 arcsec from the cD. There is also little evidence in the X-ray image for the SE pointing tail to the filaments seen in H$\\alpha+$NII and the $U$-band.  We note that the centroid of the outer X-ray emission (Fig.~4) lies not on the cD but partway (10--20 arcsec S of the cD) along the filament. This suggests that the cD is oscillating about the cluster core. There is no strong evidence from the X-ray surface brightness for a deep cusp at the cluster centre as determined from the outer emission. This argues against any static dark matter cusp at the cluster centre. The lack of any extensive ($>10$~arcsec) peak around the cD also argues against any deep dark matter well associated with the massive cD halo seen in the cluster (Johnstone, Naylor \\& Fabian 1991) from moving directly with the cD itself.  There are several possible explanations for the filament. Detailed modelling is beyond the scope of this Letter. It may a) be a cooling wake, produced by a cooling flow occurring around a moving galaxy (e.g. NGC5044; David et al 1994), b) a contrail produced by the ram pressure of the radio source passing through a multiphase medium, c) evaporation of cold gas ram-pressure stripped from the cD, or d) an accretion wake (Sakelliou, Merrifield \\& McHardy 1996).  This last possibility, d), requires that the gas be focussed by the gravitational effect of the cD and then cool into a wake. Since the likely space velocity and the velocity dispersion of the cD are both at least a factor of two smaller than the sound speed of the hot gas it is unlikely that gravitational focussing of the hot gas alone is responisble. The evaporation scenario c) requires that there was more than $5\\times 10^9\\Msun$ of gas (the hot gas mass of the filament) within the cD to be stripped out. It is not clear that it would be pulled out into such a straight, even filament (the ram-pressure-stripped wake behind the elliptical galaxy M86 in the Virgo cluster, Forman et al 1979, is very broad). Note that it is most unlikely that there is, or has been, any significant ($\\gg10^9\\Msun$) hot interstellar medium intrinsic to the cD galaxy. The radiative cooling time of any 1~keV component in the galaxy would be only $10^8\\yr$ or less and it is replenished from stellar mass loss at only a few $\\Msunpyr$ (Takeda, Nulsen \\& Fabian 1984). The cool gas at the position of the cD is then due to a combination of stellar mass loss and inflow from the intracluster medium.  We cannot rule out the contrail, b), but it is likely that the passage of the galaxy would heat gas rather than lead to it cooling. The simplest explanation is a cooling wake, a). The basic picture then is that the hot gas is in almost hydrostatic equilibrium in the cluster core potential, and cooling radiatively very slowly and steadily. The gas with a cooling time exceeding $10^9\\yr$ only senses the mean gravitational field of the cD, which then averages to a line along its direction of motion. Gas with shorter cooling time cools onto that line as the cD passes. Basically the mass of the cD serves to focus the inner part of the cooling flow, which is probably multiphase. Cooling will limit the length of the wake, in accord with the coolest X-ray gas being observed to be most distant from the cD. Of course, if the cD is also oscillating about the cluster dark matter centre then the filament length will not exceed the oscillation length.  The emission-line velocity map presented by Hu et al (1985) shows that most of the filament has the velocity of the cluster (i.e. $\\sim 150\\kmps$ blueshifted from the velocity of the cD; the cD is redshifted relative to the cluster by the same amount, Oegerle \\& Hill 1994). Only the line emission at the position of the cD shares its redshift. This velocity structure agrees with the above cooling model and argues strongly against the stripping hypothesis, where a strong velocity gradient is expected.  The extensive X-ray/H$\\alpha$ filament in A1795 offers the possibility to identify the source of energy and ionization of the optical nebulosity common in cooling flow clusters (Crawford et al 1999 and references therein; Donahue et al 2000). The filament extends at least 80 kpc from the cD and most of it is therefore well removed from the cD nucleus, the radio source, and the bulk of the stars in the cD. If part of the $U$-band emission (McNamara et al 1996a) is due to massive young stars then the formation of intergalactic stars can be studied. It will clearly be interesting to observe the clumpiness of the emission. As reported above, the optical emission from the filament does not appear to correlate in detail with the X-ray emission, which could be due to clumpiness in the gas: the denser parts cool out fastest and so form stars more readily.  The Balmer decrement of 3--4 reported by Hu et al (1985) in the filament is consistent with a small amount of reddening. Further optical spectra can search for emission lines from calcium (Donahue \\& Voit 1993; Fabian, Johnstone \\& Daines 1994) which depletes onto dust. If dust is present, it need not rule out condensation from the intracluster medium since large grains may survive sputtering and supernovae from massive stars in the filament can inject dust. This could be helpful in studying prompt dust formation from young stellar populations which is not clearly understood at present (Feigelson, private communication).  The filament also enables studies of the survivability of cool gas in the intracluster medium against mixing and conduction. Mixing layers (Begelman \\& Fabian 1990) may provide some of the photoionizing radiation for the cold gas in the filament. Thermal conduction, if it proceeds at the Spitzer value, would destroy the filament in $\\sim 10^8\\yr$, so it must be suppressed by at least an order of magnitude, presumably by magnetic fields.  In summary, the optically-derived velocity structure of the filament associates it with the intracluster gas, not the galaxy. The X-ray-derived temperature structure of the filament, with the gas furthest from the galaxy (therefore oldest) being the coolest, indicates that the gas is cooling with time and not heating up (evaporating). Together with the approximate match in timescales, the evidence points to the filament originating from cooling of the intracluster gas, attracted into a wake along the path of the moving cD galaxy."
    },
    "0011/astro-ph0011292_arXiv.txt": {
        "abstract": "s{MACHe3 (MAtrix of Cells of superfluid \\hetro) is a project of a new detector for direct Dark Matter Search. The idea is to use superfluid \\hetrois as a sensitive medium. The existing device, the superfluid \\hetrois cell, will be briefly introduced. Then a description of the MACHe3 project will be presented, in particular the background rejection and the neutralino event rate that may be achieved with such a device.} ",
        "introduction": "Many existing experiments for direct Dark Matter (DM) search \\footnote{In particular, we shall suppose all through this work a neutralino (\\neutt), the lightest supersymmetric particle, as the particle making up the bulk of galactic cold DM.} are imposing upper limits on the interaction cross-section of neutralino (\\neutt) on proton, both scalar and axial. They use a great variety of materials and detection techniques. However, the neutrons are still very difficult to discriminate along with a good rejection for \\gamss. The cosmogonic activation is, in all these experiments, an important source of intrinsic contamination. There are several advantages of using \\hetrois as a sensitive medium for direct DM detection. Some of them are : \\begin{itemize} \\item It presents a high neutron capture cross-section (up to $\\sim \\!10^3$ barn for $\\sim \\!0.1$ eV thermal neutrons), giving a clear neutron signature and hence a good rejection, as shown in sec \\ref{sec:back}. As neutrons interact {\\it a priori} like \\neutt, this gives to an \\hetrois detector the possibility to discriminate the \\neut signal from its ultimate background noise. \\item Superfluid \\hetrois is produced with a very high purity, making this medium almost free of radioactive contaminations. \\item The interaction of a \\neut on the \\hetrois nucleus (odd spin) is mainly a spin-dependent one, making this medium complementary to other media used in existing DM experiments. \\end{itemize} In addition a DM detector should present an energy threshold as low as possible. At ultra low temperatures (T $\\simeq \\!100 \\,\\mu\\mathrm{K}$), \\hetrois in the superfluid phase B presents in a Lancaster's cell configuration \\cite{lanc} a low detection threshold ($\\mathrm{E}_{th}\\simeq 1 \\,\\mathrm{keV}$).\\\\ In order to enhance the rejection against the background (neutrons, \\gams and muons), a high granularity superfluid \\hetrois detector (fig.~\\ref{fig:mat}) as been proposed~\\cite{firstmac3,dm20}. This matrix configuration would allow to obtain a high background rejection (sec~\\ref{sec:back}) along with a reasonable \\neut event rate (sec.~\\ref{sec:susy}). ",
        "conclusions": "MACHe3 would allow to obtain a high intrinsic rejection against $\\gamma$-ray, neutron and muon background, by means of correlation among the cells and energy loss measurement. The contamination is evaluated, for $\\mu$ and neutrons, as well below the maximum expected \\neut rate. The results of the simulation are encouraging, both for signal and background rejection."
    },
    "0011/hep-ph0011134_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ \\noindent We present numerical simulations of fragmentation of the Affleck-Dine condensate in two spatial dimensions. We argue analytically that the final state should consist of both Q-balls and anti-Q-balls in a state of maximum entropy, with most of the balls small and relativistic. Such a behaviour is found in simulations on a 100$\\times$100 lattice with cosmologically realistic parameter values. During fragmentation process, we observe filament-like texture in the spatial distribution of charge. The total charge in Q-balls is found to be almost equal to the charge in anti-Q-balls and typically orders of magnitude larger than charge asymmetry. Analytical considerations indicate that, apart from geometrical factors, the results of the simulated two dimensional case should apply also to the fully realistic three dimensional case. ",
        "introduction": "The Minimal Supersymmetric Standard Model (MSSM) has several flat directions where the scalar potential is nearly identically zero \\cite{dinert}. During inflation, squark and slepton fields will fluctuate freely along the flat directions, forming Affleck-Dine (AD) -condensates  \\cite{ad}. The  state of lowest energy, however, is not the AD-condensate but a non-topological soliton, the Q-ball \\cite{cole2,ks1}, which carries a non-zero baryonic (B-ball) or leptonic (L-ball) charge. The instability is induced by the spatial perturbations \\cite{ks,bbb1} which are naturally present in the condensate because of quantum fluctuations during inflation. The fragmentation process and the properties of the resulting Q-balls depend on SUSY breaking: if gauge mediated, Q-balls will be large and completely stable \\cite{ks,ksd} and form along every flat direction\\footnote{For another variant of gauge mediated Q-balls, see \\cite{kasuya3}.} and could be detectable even today \\cite{kkst}, whereas if SUSY breaking is gravity mediated, Q-balls will be unstable  \\cite{bbb1,bbb2} but long-lived enough to have a host of interesting cosmological consequences \\cite{bbb2,bbbdm}. In the gravity mediated case the formation of Q-balls is  also a generic feature in all but few flat directions \\cite{ejm}.  Fragmentation of the AD-condensate involves highly non-linear dynamics and is therefore quite complicated \\cite{adfrag,mv1}. The condensate lumps must somehow loose the extra energy when settling down to a Q-ball configuration. This they may do by radiating out quanta of AD-scalars or smaller condensate lumps, but the lumps  may also experience frequent collisions. Since for a fixed charge the energy of a Q-ball is also fixed, fragmentation dynamics and the final distribution of Q-balls will obviously depend on the initial energy and charge density of the condensate. In principle, this is a free parameter, although the observed baryon asymmetry can be argued to indicate a natural order of magnitude for the ratio.  In Ref. \\cite{adfrag} the relaxation of a single spherical condensate lump was followed numerically, but sphericality is likely to be a far too constraining assumption. Moreover, the condensate fragments are not isolated but interact with each other. The general features of Q-ball collisions have been studied in Refs. \\cite{collisions,mv2,mv3}; the cross section in the gravity mediated case has been shown to depend on the relative phases of the colliding Q-balls \\cite{mv2}. Q-ball formation has been seen in 3-d lattice studies both in gauge mediated \\cite{kasuya1} and gravity mediated \\cite{kasuya2} cases but only for a single energy-to-charge ratio of the condensate in each case. As we will show in this paper, natural choices for the initial energy-to-charge ratio necessarily lead to a copious formation of not only Q-balls but also of anti-Q-balls, which carry a net negative charge. Such a conclusion can be reached already by relatively simple analytical arguments which assume that the condensate lump reaction rates are fast enough to drive them into a state of maximum entropy. From the resulting equilibrium distributions it then follows that the number of Q-balls is almost equal to the number of anti-Q-balls so that the total positive (negative) charge in Q-balls is typically much larger than the actual charge asymmetry. Indeed, as we shall discuss in this paper, detailed numerical simulations show this behaviour.  The paper is organized as follows. In Section 2 we derive the Q-ball and anti-Q-ball number density distributions by assuming that the condensate lumps actually equilibrate; this is shown to be a self-consistent assumption. We also show that as far as the distributions are concerned, there is no qualitative difference between two and three spatial dimensions. In Section 3 we present the results of numerical simulations of the condensate fragmentation in 2+1 dimensions. We follow the time evolution of the condensate field, show it breaking up into Q-balls and anti-Q-balls, and extract the number density distributions from the numerical data. In Section 4 we present our conclusions. ",
        "conclusions": "We have studied the process of Q-ball formation from the AD-condensate in the early universe by numerical simulations in two spatial dimensions. The analytical considerations presented in Sect. 2 reveal that a typical Q-ball is relativistic. This conclusion holds if at least part of the baryon asymmetry of the universe is contained in baryon number carrying Q-balls. The reaction rate of Q-ball interactions is then much larger than the Hubble rate and hence we may expect that a state of maximum entropy is likely to be reached. Assuming a thermal Q-ball distribution we showed that in a typical cosmological situation the AD-condensate fragments into practically equal numbers of Q-balls and anti-Q-balls so that $N_+\\approx N_-$. The  amount of charge contained in these balls is much larger than the net charge, $Q_+\\approx Q_-\\gg Q_{{tot}}$. Estimation of the average velocity of a Q-ball confirms that Q-balls are typically relativistic. The analytical calculations were carried out in both two and three dimensions and we found that the shapes of the number density distributions or velocities are only very weakly dependent on dimension. These considerations strongly support the assumption that 2-d simulations can in fact capture the essential features of the real 3-d case.  Our numerical simulations were carried out on a $2+1$ dimensional $100\\times 100$-lattice. The initial condition was chosen to be uniform AD-condensate with small perturbations. The initial energy-to-charge ratio $x$ was varied from $1$ to $\\sim 10^5$. We found that regardless of $x$ the qualitative behavior of the AD-condensate is similar in the beginning of the simulation. The small perturbations start to grow and due a complex dynamical process the AD-condensate fragments into several lumps carrying positive charge. These are not yet Q-balls but some excited states \\cite{adfrag}.  {}From this point on the further evolution of the condensate depends on the actual value of $x$. If $x=1$, the lumps slowly relax into Q-balls and no anti-Q-balls are present. This is what we would expect from analytical studies and this is what has been seen in 3-d lattice simulations \\cite{kasuya1,kasuya2}. If, however, $x\\gg 1$, the condensate fragments into a large number of Q- and anti-Q-balls, as we have shown here. We also showed that the distributions of the Q- and anti-Q-balls obtained from the simulations are thermal, which again is in accordance with analytical expectations.  An interesting feature which is apparent in Figs. \\ref{form7}-\\ref{form9} is that the system of AD lumps appears to possess some non-trivial spatial structure. First, the large Q- and anti-Q-balls can be seen to form in pairs which give rise to a positive spatial correlation between them. It should be noted that if small balls are also formed in pairs they are likely to have a homogenized spatial distribution due to their large velocities. Second,  there are long range correlations which show up as filament-like structures in Figs. \\ref{form6}-\\ref{form8}. The texture visible in the middle phase of the non-linear evolution may well survive in the spatial Q-ball distribution also at later times, although it no longer is visible. Numerical correlation analysis indeed shows that still at the final stage of our simulation (Fig. \\ref{form9}) the mean distance between both Q- and anti-Q-balls deviates by over four standard deviations from the random uniform distribution. The deviation is negative \\ie the mean distance between Q-balls is smaller than expected for a random distribution. These features may affect observables such as the cosmic microwave background or large scale structure and deserve further study.  The fact that the Q-ball distribution in the cosmologically realistic case, $x\\gg 1$, tends toward a thermal one implies, like the analytical treatment shows, that a large number of Q-balls move relativistically. The very large Q-balls that are clearly visible in Fig. \\ref{form12} appear immobile but there is also a large number of smaller fast moving Q-balls which can thermalize the distribution. In a thermodynamical sense one can view the fragmentation of the AD-condensate into Q-balls as a relaxation process where entropy is maximized: from this point of view the appearance of the thermal distribution is natural.  In this work we left the gauge mediated SUSY breaking scenario without the attention it deserves. This is due to a number of facts that make the simulation of the gauge mediated SUSY breaking scenario more cumbersome compared to the gravity mediated case. In the fragmentation process the spectra of growing modes play an important role since the first mode that goes non-linear determines the size of the fastest growing disturbance in the condensate. The SUSY breaking mechanism, \\ie whether it is gauge or gravity mediated, determines the size of the fastest growing mode via the actual form of the scalar potential. As we have seen, the size of the instabilities in the gauge mediated case are larger than in the gravity mediated scenario. Therefore one needs a lattice with a much larger physical size. Obviously, one can always increase lattice spacing but this happens at the cost of losing resolution. Moreover, the large instabilities correspond to large Q-balls forming from the AD-condensate which in the gauge mediated scenario have thin walls. To avoid numerical instabilities due to the thin walls a high resolution is required. Therefore, in order to avoid finite size effects, one would need a much larger number of lattice points than in the gravity mediated case.  In addition to the numerical difficulties the gauge mediated scenario does not benefit form the analytical arguments presented in Sect. 2. This is due to the different energy-charge -relations of the two cases. For large Q-balls in the gauge mediated case the energy is related to the charge by $E\\sim Q^{3\\over 4}$ instead of the linear relation $E\\sim Q$ in the gravity mediated scenario. From (\\ref{part1}) one can see that the Q-integral diverges since the $\\exp (-\\beta E+\\mu Q)$-term is dominated by $\\mu Q$ in the large Q limit in the gauge mediated scenario. Physically the lack of a proper asymptotic equilibrium state implies that the system tends towards a state where charge is concentrated in large balls (and ultimately in a single Q-ball) as noted also in Ref.\\ \\cite{ls}. However, one should also take into account that the more charge is stored in large balls the more there exists leftover energy in the system that must be accounted for. The final state of the system would then be strongly dependent on the time at which the distribution freezes due to the expansion of the universe and hence on the initial conditions.  To conclude, our results indicate that in the gravity mediated SUSY breaking case, the AD condensate fragments into an almost equal number of  Q-balls and anti-Q-balls. The evaporation times of the Q-balls, the nature of particles they decay into, and possible spatial correlations surviving the fragmentation process and the subsequent decays of Q-balls, are obviously very important factors that determine the signatures left by a possible Q-ball period in the early history of the universe."
    },
    "0011/astro-ph0011151_arXiv.txt": {
        "abstract": "We show the results of two--dimensional, azimuthally symmetric simulations of advection dominated fluid flows around black holes. We use the SPH method, and the $\\alpha$--prescription for the viscosity, including all terms in the stress tensor. We have performed calculations with $0.001 < \\alpha < 0.2$ and find strong circulation patterns in the flow, on spatial scales that depend on the magnitude of the viscosity. ",
        "introduction": "The astrophysical relevance of accretion flows in which radiative cooling is inefficient has been recently emphasized~\\cite{NY94}. These flows show complicated hydrodynamical behavior that is very different from that obtained with the thin disk approach~\\cite{SS73} for efficient cooling.  We use the Smooth Particle Hydrodynamics (SPH) method~\\cite{Monaghan}, in azimuthal symmetry. The equation of state is $P=\\rho u (\\gamma-1)$ with $\\gamma=4/3$. Viscosity is computed using an $\\alpha$ prescription ($\\mu=\\alpha \\rho c_{s}^{2}/\\Omega_{k}$), including all the terms from the viscous stress tensor. The black hole is modeled as a Newtonian point mass with an absorbing boundary at $r_{in}=1.3r_{G}=1.3(2GM_{BH}/c^{2})$ . We use two kinds of initial conditions: 1) equilibrium hydrostatic tori with constant specific angular momentum; 2) matter injection at a steady rate at a radius $r=300r_{G}$. We assume no radiative energy losses. ",
        "conclusions": ""
    },
    "0011/astro-ph0011367_arXiv.txt": {
        "abstract": "The solar tachocline, located at the interface between the latitude-dependent rotation of the convection zone and the rigid radiative interior, presents high gradients of angular velocity which are of particular interest for the models of the solar dynamo and angular momentum transport. Furthermore, latitudinal and temporal variations of the tachocline parameters, if any, are also of particular interest in order to constrain models. We present a review of some of the theories of the tachocline and their predictions that may be tested by helioseismology. We describethe methods for inferring the tachocline parameters from observations and the associated difficulties. A review of results previously obtained is given and  an analysis of the new 6 years database of LOWL observations is presented which yields no compelling evidence of variations or general trend of the tachocline parameters during the ascending phase of the current solar cycle (1994-2000). ",
        "introduction": "\\label{sec:introduction} The solar tachocline, so called after \\citet{spiegel92},  is defined as the layer at the base of the convection zone (hereafter, CZ) where important radial  shear occurs. It is the transition zone where the angular velocity changes from its latitude-dependent value in the CZ to its constant and intermediate value in the radiative interior. We can emphasize three main reasons why this layer is of particular interest.  {\\bf(i)}Shear turbulence and/or meridional circulation inside the tachocline may provide a mechanism for mixing material between the CZ and the radiative interior \\citep[e.g.][]{brun99,schatzman00} which is needed to understand, for instance, the burning of lithium and the depletion of helium and to reach a better agreement between solar models and helioseismic observations \\citep[e.g.][]{richard96,brun99}.  {\\bf(ii)}The tachocline may be the seat for angular momentum transport processes that could lead to the observed rigid rotation rate of the radiative interior. Hydrodynamical transport by unstable shear flow \\citep[e.g.][]{chaboyer95} and transport by internal gravity waves in the tachocline \\citep{kumar_talon99,kim00} have been studied but they have not been found efficient enough to lead to an uniform internal rotation  at the present age of the Sun. This probably requires also the presence of an internal magnetic field \\citep{mestel_weiss,Charbonneau_mcgregor93,gough_mcintyre}.  {\\bf (iii)} Finally, the tachocline is  the best location for an oscillatory solar dynamo which is generaly  believed to be responsible for the solar magnetic cycle for the following reasons: (1) Its  radial and latitudinal differential rotation  has the ability to produce a toroidal field by shearing a pre-existing poloidal field. (2) The $\\alpha$-effect, the essential mechanism for producing poloidal field from toroidal field, is usually located in the CZ but the tachocline can also produce a strong $\\alpha$-effect  by magnetic buoyancy instability \\citep{ferriz94} and/or by the unstable shallow-water modes \\citep{dikpati_gilman_00_1}. (3) Because the tachocline (or  part of it) may also be located in the slight sub-adiabatic overshoot layer, the toroidal fields can be stored  for an extended period of time, and therefore can be amplified and acted on by the $\\alpha$-effect before they escape to the surface through buoyant rise or be disrupted completely by convective shredding.      We will first (Sect.~\\ref{sec:layers}) give some terminology about the different layers and their properties at the base of the CZ (also to be defined). Then, in Sect.~\\ref{sec:models}, we present some models of the tachocline and their predictions that may be tested by helioseismology using methods presented in Sect.~\\ref{sec:methods}. Previous results obtained are summarized in Sect~\\ref{sec:eq} and we then  use the new LOWL data to investigate latitudinal and temporal variations of the tachocline parameters during the  ascending phase of the current solar cycle (Sect.~\\ref{sec:variations}). ",
        "conclusions": ""
    },
    "0011/astro-ph0011221_arXiv.txt": {
        "abstract": "The most recent experimental results on the cosmic ray composition and energy spectrum above 1 TeV are reviewed and discussed. All data agree on the presence of the so-called ``knee'' at an energy $E_{k} \\simeq 3 \\ 10^{15} eV$; the knee is seen in all the components of the Extensive Air Showers. These results support the hypothesis of an astrophysical origin of the knee, while no new features  in the hadronic interactions at high energies are envisaged.\\\\ The cosmic ray composition below and above the knee region is still an open question. According to most experiments, the knee seems to be due to the light component of the primary beam, with a composition getting heavier above the knee. However, results contradicting this conclusion have to be considered and understood. ",
        "introduction": "The energy spectrum of cosmic rays (CR) spans a very wide energy range, with particle fluxes steeply falling more than 30 orders of magnitude. Above the solar modulation region, the spectrum can be well described by a power law, which steepens around $3 \\ 10^{15} eV$, a feature called the ``knee'', discovered in 1958 \\cite{kri}; it softens again at $\\simeq 10^{19} eV$, the ``ankle''.\\\\ Explaining the knee feature would shed light on the CR origin and acceleration mechanism, depending on whether it is a signature of a change in the hadronic interactions at such energies or it reflects a feature of the cosmic ray spectrum, thus concerning mainly astrophysics.\\\\ Several arguments involving energetics, composition and secondary $\\gamma$ ray production suggest that cosmic rays at least up to the knee region are confined in the Galaxy. \\\\ The most popular theory is that of diffusive shock acceleration in Supernova remnants (SNR), that is particle acceleration by SNRs expanding supersonically in the surrounding medium. \\\\ Supernova explosions can easily account for the energy stored in galactic cosmic rays; the spectrum emerging from the SNR is of the type $E^{-2.1}$ up to a maximum energy near $10^{14}$ eV times the nuclear charge, after which it drops very rapidly. Folding the production spectrum with the effect of diffusion through the Galaxy, and taking a trapping time varying as $E^{-0.6}$ (as found from the proportion of secondary to primary nuclei arriving to Earth), the resulting flux of CR in the Galaxy would be $\\propto E^{-2.7}$, in close agreement with expectations. The maximum achievable energy is close to the knee one.\\\\ A direct evidence that the nucleonic component of CR is indeed produced in SNRs could be obtained by the observation of $\\gamma$ rays: the accelerated cosmic rays can in fact interact with the local interstellar matter, in this way producing $\\gamma$ rays by either hadronic or leptonic production.\\\\ Various experimental groups reported on TeV $\\gamma$ emission from supernova remnants, like SN1006, RJX1713.7-3946, Cassiopea-A (see e.g. \\cite{weekes} and references therein). Unfortunately in all cases the emission can be attributed to electron progenitors and no positive evidence for hadroproduction of TeV $\\gamma$'s has been found yet.\\\\ Further information can be obtained from the study of the distribution of CR arrival directions. A recent compilation of the anisotropy measurements can be found in \\cite{anis}; while the amplitude and phase of anisotropy data below $\\simeq 2 \\ 10^{14} eV$ are consistent and statistically accurate \\cite{et-ani}, the experimental results at higher energies are not compatible with expectations. This could mean that the diffusion model cannot be simply extrapolated to these energies \\cite{hillas}.\\\\ Various models have been put forward trying to identify the sites and mechanisms of injection of cosmic rays at higher energies, at and above the knee. If the bend in energy spectrum is related to the maximum achievable energy in the accelerator, then CR at higher energies could be powered by a reacceleration by interstellar turbulence \\cite{Seoptu}, or they could be produced by Supernovae exploding in denser media (their stellar wind cavity) (\\cite{Bie98} and references therein). On the other hand, the knee could be attributed to propagation effects. In both cases, one would expect multiple bends due to the different elements bending at fixed rigidity; the composition would become heavier above the knee. An extra-galactic origin for CR above the knee has also been proposed \\cite{prot}, where the accelerator sites are found in Active Galactic Nuclei and the resulting composition is getting lighter above the knee.\\\\ A completely different point of view assigns the knee to a new dramatic process of hadronic interaction which takes over around the knee energy. However, even if it is true that in the $\\Delta E$ of interest we have no direct information about the hadronic interaction cross section for the secondary production relevant to the interpretation of measurements, no experimental data as far show a need for a different interaction mechanism.\\\\ From the experimental point of view, what is most important in order to test the models is to measure the cosmic ray composition and energy spectrum near the energy limit of the shock models; moreover, measurements of anisotropy and secondary to primary ratio at higher energy are of utmost importance. ",
        "conclusions": "The cosmic ray energy spectrum and composition are studied with a variety of experimental techniques detecting different air shower components in the energy region above 1 TeV.\\\\ Below the knee, no new data are available and the conclusions reached by JACEE and RUNJOB experiments still hold; however new projects, planned to fly on balloons or on the Space Station, are in progress. They will surely extend the explorable energy region and the available statistics on single nuclei.\\\\ The energy spectrum has been studied in detail both by charged particles and $\\breve{C}$erenkov light ground arrays and some firm conclusions were reached: all data agree on the existence of the knee in the primary energy spectrum of cosmic rays at an energy $\\simeq 3-4 \\ PeV$. The bend has been seen in all shower components, thus supporting an astrophysical interpretation of the knee as opposite to that of a change in the hadronic interaction picture at these energies. All results agree to attribute the knee to the medium-light mass primaries.\\\\ Progress has been made as regards the mass composition. Almost all the ground array results show an increase in the primary mean logarithmic mass above the knee, even if the absolute scale can be quite different. There are however contradicting results coming from experiments relying on the $\\breve{C}$erenkov light detection from Air Showers; the differences in the measure of $X_{max}$ are however quite big.\\\\ It is very important to study in detail the systematics which could bias the results; the combined use of different observables and the comparison among various data sets can help in this task. The different sensitivity to composition of the various used observables, the methods employed to determine the primary energy and the problems found in the models used in the simulations could explain the spread in the results which is apparent in Fig.\\ref{fi:compall}. \\\\  {\\bf Acknowledgements}  I would like to thank my friends G.Navarra, B.Alessandro and A.Chiavassa for the interesting and useful discussions which were the basis of this paper. Sincere thanks also to all my Brazilian friends, for their warm hospitality during my stay in Campinas."
    },
    "0011/nucl-th0011003_arXiv.txt": {
        "abstract": "We discuss the kaon-nucleon interaction and its consequences for the change of the properties of the kaon in the medium. The onset of kaon condensation in neutron stars under various scenarios as well its effects for neutron star properties are reviewed. ",
        "introduction": "\\label{sec:intro}  Due to its non-abelian structure, Quantum Chromodynamics (QCD) becomes very strongly interacting and highly nonlinear at large space-time distances. As a consequence, quarks and gluons condense in the physical vacuum with a gain in condensation energy density of $\\Delta\\epsilon_0\\sim 500$ MeV/fm$^3$. The condensation of quarks is associated with the spontaneous breaking of chiral symmetry, an additional symmetry of the QCD Lagrangian in the absence of (current) quark masses. This limit is well justified in the up-down quark sector and to a somewhat lesser extent for the strange quark. There is good evidence that the mechanism for spontaneous chiral symmetry breaking is provided by classical gluon field configurations in euclidean space called 'instantons'. These provide effective quark-(anti)quark interactions which are strong enough to cause a BCS like transition to a condensed state of quarks and antiquarks. It has been shown, that this picture provides an excellent description of hadronic states and correlation functions \\cite{SchaSh} and it is fair to say that the low-lying hadron spectrum is dominated by spontaneous chiral symmetry breaking with confinement playing a much lesser role. These observations form the basis for discussing the properties of hadrons, or more precisely hadronic correlation functions, under extreme conditions in temperature and/or density as encountered in the early universe or in the interior of neutron stars. It is well established from lattice QCD that the vacuum undergoes a phase transition (or at least a sharp cross over) when heated. Chiral symmetry is restored in this phase accompanied by a nearly vanishing chiral quark condensate. Though not calculable at present from first principles, the same is expected to happen at finite density. Since light hadrons are dynamically driven by chiral symmetry breaking it is obvious to ask how hadronic properties are related to the vacuum structure and its changes with temperature and density. This is far from trivial and under intense experimental and theoretical scrutiny at present.  As detailed below the most economical way for treating hadrons in matter under extreme conditions is to resort to 'effective field theories' in which hadrons rather than quarks and gluons appear as the fundamental degrees of freedom. Formally one identifies the pertinent quark currents $J_{\\Gamma_j}(x)=\\bar q(x)\\Gamma_j q(x), \\Gamma_j=1, \\gamma_5,\\gamma_\\mu,..$ with elementary hadronic fields $\\phi_i(x)$ for which the most general effective Lagrangian, consistent with the underlying symmetries and anomaly structure of QCD, is written down. We recall that the spontaneous breaking of chiral symmetry has two important consequences. One is the appearance of (nearly) massless Goldstone bosons (pions, kaons, eta) and the other the absence of parity doublets in the hadron spectrum ($m_\\pi\\neq m_{f_0}, m_\\rho\\neq m_{a_1}$ etc). For the present discussion the first is the most relevant. Chiral symmetry does more than just predict the existence of Goldstone bosons. It also prescribes and severely restricts their mutual interactions as well as those with other hadrons. The most rigorous treatment is in terms of 'chiral perturbation theory' \\cite{GaLe} but other 'chiral effective Lagrangians' including e.g.\\ vector mesons can be devised \\cite{NK,NK97}.  To elucidate the connection between the vacuum structure (the chiral condensate) and the properties of light hadrons let us consider a medium of hadronic matter in thermal and chemical equilibrium. This is of course well suited for neutron stars. The QCD partition function is then given in the grand-canonical ensemble as \\beq {\\cal Z}_{QCD}(V,T,\\mu_q)={\\rm Tr}\\exp{\\{-(\\hat H_{QCD}-\\mu_q\\hat N_q)/T\\} } \\quad , \\eeq where $\\hat H_{QCD}$ denotes the QCD Hamiltonian, $\\hat N_q$ the quark number operator and $\\mu_q$ the quark chemical potential. Statistical expectation of operators are then given as \\beq \\tave{\\hat O}={\\cal Z}^{-1}\\sum_n\\bra{n}\\hat O\\ket{n}\\exp{\\{-(E_n-\\mu_q N_n)/T\\} } \\quad , \\eeq where $E_n$ are the exact QCD energies (hadrons). The quark condensate $\\tave{\\bar qq}$ can be obtained directly from the free energy density \\beq \\Omega_{QCD}(T,V)=-\\lim_{V\\to\\infty}{T\\over V}\\ln{\\cal Z}_{QCD}(V,T,\\mu_q) \\eeq via the Feynman-Hellmann theorem as \\beq \\tave{\\bar qq}={\\partial\\Omega_{QCD}\\over\\partial m^\\circ_q} \\quad , \\eeq where $m^\\circ_q$ denotes the bare (current) quark mass. An obvious first step is to approximate the free energy density by an ideal gas of hadrons. Using the Gell-Mann Oakes Renner relation for the vacuum chiral condensate \\beq m_\\pi^2f_\\pi^2=-2m^\\circ_q\\ave{\\bar qq} \\quad , \\eeq where $m_\\pi$ is the pion mass and $f_\\pi$ the pion weak-decay constant, one then finds \\beq {\\tave{\\bar qq}\\over\\ave{\\bar qq}}=1- \\sum_h{\\Sigma_h\\varrho^s_h(\\mu_q,T)\\over f_\\pi^2 m_\\pi^2} \\quad. \\label{qqt} \\eeq Here \\beq \\Sigma_h=m^\\circ_q{\\partial m_h\\over \\partial m^\\circ_h} \\eeq denotes the 'sigma' commutator' (related to the scalar density of quarks in a given hadron) and $m_h$ the vacuum mass of a given hadron. At low temperature and small baryochemical potential ($\\mu_B=3\\mu_q$), in which the hadron gas is dominated by thermally excited pions and a free Fermi gas of nucleons,  (\\ref{qqt}) leads to the celebrated leading-order result \\beq {\\tave{\\bar qq}\\over\\ave{\\bar qq}}=1-{T^2\\over 8f_\\pi^2} -0.3{\\rho\\over\\rho_0} \\dots \\eeq where $\\rho_0=0.16$/fm$^3$ is the saturation density of symmetric nuclear matter. This result is model independent and indicates that the mere presence of an ideal gas of hadrons already alters the vacuum structure and leads to a decrease of the condensate, without changing the vacuum properties of the hadrons! Obviously medium-modifications of hadrons and the corresponding non-trivial change of the QCD vacuum has to involve hadronic interactions. They become increasingly important as the medium grows hotter and denser, i.e.\\ as the point of chiral restoration is approached. Thus, the theoretical description in terms of hadrons becomes very complex, involving more and more degrees of freedom. In the vicinity of the restoration transition, in addition, non-perturbative methods are called for which is far from trivial in effective field theories.  In terms of effective fields $\\phi_j(x)$ representing the pertinent quark currents $J_{\\Gamma_j}(x)$, hadronic correlation functions in a hot and dense environment are defined as the (retarded) current-current correlation functions \\beq D_{\\phi_j}(\\omega,\\vec q)=-i\\int d^4x\\theta(x_0)\\tave{[\\phi_j(x),\\phi_j(0)]}. \\label{prop} \\eeq Note that, in contrast to the vacuum, these correlators depend on energy $\\omega$ and three-momentum $\\vec q$ separately since Lorentz invariance is explicitly broken by the presence of matter. Equation (\\ref{prop}) can be rewritten in terms of the self energy $\\Sigma_{\\phi_j}$ as \\beq D_{\\phi_j}(\\omega,\\vec q)=(\\omega^2-\\vec q^2- m_{\\phi_j}^2-\\Sigma_{\\phi_j}(\\omega,\\vec q))^{-1} \\eeq where $m_{\\phi_j}$ denotes the (bare) mass of the field $\\phi_j$ and all interaction effects are incorporated via $\\Sigma_{\\phi_j}$ which depends on $T$ and $\\mu_B$. Given the effective Lagrangian, the objective is then to evaluate $\\Sigma_{\\phi_j}$ as realistic as possible. This is usually attempted by employing 'chiral counting rules' for evaluating loop diagrams contributing to $\\Sigma_{\\phi_j}$ in the low-density and low-temperature limit and by adjusting the parameters of the effective Lagrangian to as many data for elementary scattering processes as available. If everything is done consistently, chiral symmetry is properly incorporated and the relation between hadronic medium modifications and vacuum changes can be inferred.  In the vacuum, the hadronic correlators $D_{\\phi_j}$ are usually dominated by a few fairly sharp hadronic 'resonances'. These are visible as 'peaks' in the corresponding 'spectral functions' \\beq \\rho_{\\phi_j}(\\omega,\\vec q)=-{1\\over\\pi}{\\rm Im}D_{\\phi_j}(\\omega,\\vec q) \\eeq at the hadronic vacuum mass $m_{\\phi_j}$. Two things will happen in the interacting medium. One the one hand the mass will change to an effective 'pole-mass' \\beq m_{\\phi_j}^{*2}=m_{\\phi_j}^2+{\\rm Re}\\Sigma_{\\phi_j}(m_{\\phi_j}^{*2},0) \\quad . \\eeq In fact, one may have \\emph{several} solutions and the in-medium spectrum shows more than one 'peak'. As we shall see, this is usually the case. On the other hand, the imaginary part ${\\rm Im}\\Sigma_{\\phi_j}$ acquires additional pieces through the interactions with the medium giving rise to an increased width. If the width becomes too large, the peak structure is washed out and the notion of a 'quasiparticle' is lost. This happens when the 'quasiparticle' energy \\beq \\omega_{\\phi_j}^2(\\vec q)=\\vec q^2+m_{\\phi_j}^2+{\\rm Re} \\Sigma_{\\phi_j}(\\omega_{\\phi_j}^2,\\vec q) \\eeq is no longer large compared to the width \\beq \\Gamma_{\\phi_j}(\\omega_{\\phi_j}(\\vec q))= -{1\\over 2\\pi}{\\rm Im}\\Sigma(\\omega_{\\phi_j},\\vec q) \\quad. \\eeq This has to be kept in mind when describing in-medium hadrons.  In the following, we will discuss the properties of kaons in the medium and its consequences for the properties of neutron stars. In sect.~\\ref{sec:AR}, the elementary kaon-nucleon interaction and the in-medium changes of the kaons as seen in kaonic atoms are described. Chiral effective interactions are constructed which can describe these data and are used to extract the optical potential of kaons at finite density. Section~\\ref{sec:JSB} utilizes the results found for the kaon optical potential to apply it to the equation of state for neutron stars. Implications for a first order phase transition to a kaon condensed state are listed. Effects on the onset of kaon condensation by the presence of hyperons in matter are studied. Finally, we summarize in sect.~\\ref{sec:summary} and give an outlook. ",
        "conclusions": "\\label{sec:summary}  We have discussed the elementary kaon-nucleon interaction as derived from scattering data and the in-medium changes of the K$^-$ as deduced from kaonic atoms. Various coupled channel calculations using chiral effective interactions or the boson exchange model with higher order partial waves and their results for the K$^-$ properties in dense matter have been reviewed. The extracted range of the optical potential has been used to study the phase transition to kaon condensation in neutron star matter. We outlined the effects of a first order phase transition due to the presence of structures in the mixed phase. Consequences for the global properties of neutron stars, i.e.\\ a reduced maximum mass and minimum radius, have been addressed. Finally, we examined the r\\^ole of hyperons on the onset of kaon condensation in neutron stars.  It is clear from our review, that the discussion about kaon condensation is far from being complete at present. The topic is a challenge both for experimentalists as well as for theorists, demonstrating a growing interplay between traditional nuclear physics, heavy-ion physics, and astrophysics. For example, the importance of p-wave interactions also for cold, dense neutron star matter has been stressed, but further work is needed in that direction. Especially, the precise value of the optical potential of the K$^-$ is a crucial ingredience for the neutron star matter calculation and needs to be pinned down more precisely, be it by the experimental study of deep lying levels in kaonic atoms, subthreshold kaon production in heavy-ion collisions at GSI, or by the mass-radius measurement of neutron stars.  The inclusion of hyperons for the equation of state is certainly a necessary one, but there are only a few works which have been devoted to this issue. In particular, the hyperon-hyperon interaction as well as the kaon-hyperon interaction can be important for the onset of kaon condensation and could be addressed by e.g.\\ a chiral SU(3) symmetric model.  A consistent calculation, incorporating realistic kaon-baryon as well as a baryon-baryon interactions on the same basis is still missing.  Last but not least, we point out, that kaon condensation can have impacts on other facets of neutron stars, like the evolution of proto-neutron stars and the deconfinement phase transition to quark matter. Concerning the latter point, the appearance of both, a kaon condensed phase and a quark matter phase, will lead to a triple point of strongly interacting matter inside a neutron star. At this point, the three phases, the normal hadronic, the kaon condensed and the quark matter phase, are in equilibrium. Beyond that point, the kaon condensed phase and the quark matter phase are forming a mixed phase. The presence of this triple point might have interesting implications for transport phenomena in neutron stars."
    },
    "0011/gr-qc0011093_arXiv.txt": {
        "abstract": "{ We present a new survey of the radial oscillation modes of neutron stars. This study complements and corrects earlier studies of radial oscillations. We present an extensive list of frequencies for the most common equations of state and some more recent ones. In order to check the accuracy, we use two different numerical schemes which yield the same results.  The stimulation for this work comes from the need of the groups that evolve the full nonlinear Einstein equation to have reliable results from perturbation theory for comparison. ",
        "introduction": "\\label{sec:intro} As they are the simplest oscillation modes of neutron stars, radial modes have been the first under investigation, more than 35 years ago (Chandrasekhar 1964a, 1964b). More important, they can give information about the stability of the stellar model under consideration. Since radial oscillations do not couple to gravitational waves, the appropriate equations are quite simple, and it is relatively easy to numerically solve the eigenvalue problem that leads to the discrete set of oscillation frequencies of a neutron star. In the absence of any dissipative processes, the oscillation spectrum of a stable stellar model forms a complete set; it is therefore possible to describe any arbitrary periodic radial motion of a neutron star as a superposition of its various eigenmodes.  The radial modes of neutron stars have been thoroughly investigated by various authors mostly for zero temperature equations of state (EOS) (e.g. \\cite{Har65}, \\cite{Chan77}, \\cite{GL83}, \\cite{VC92} and references therein).  But also protoneutron stars with a finite temperature EOS (\\cite{GHZ97}) and strange stars were studied (\\cite{BH91}, \\cite{VC92}, \\cite{GZ99}).  The first exhaustive compilation of radial modes for various zero temperature EOS was presented by \\cite{GL83} (hereafter GL). However, as was later pointed out by \\cite{VC92} (hereafter VC), their numerical values for the oscillation frequencies seemed to be flawed although their equations were correct. VC computed the radial frequencies for 6 equations of state of dense matter and corroborated their own results using the argument (\\cite{Har65}) that for the numerical code to be correct, it must yield a zero-frequency mode at exactly that central density for which the neutron star reaches its maximal mass. This is the point where the stellar model becomes unstable with respect to radial collapse if the central density is further increased. Yet, this is not the case for the results of GL, as was noticed by VC. However, the above mentioned test can be used only in the case when both in the stellar model and in the perturbation equations the equilibrium adiabatic index is used. In general, different adiabatic indices can be used depending on the physical conditions inside a star (\\cite{GHZ97}). For example, if the slowness of weak interaction processes are taken into account, the regions of configurations stable with respect to radial perturbations extend beyond the central density of the star with the minimum mass (e.g.~\\cite{Chan77}) and of the star with maximum mass (\\cite{GHG95}).  In this paper, we repeat the numerical calculation of the radial oscillation modes of neutron stars for various zero temperature equations of state using the equilibrium adiabatic index. To verify the results, we use two different formulations of the equations together with two different numerical methods to solve the eigenvalue problem. We find that in all cases we obtain matching values for the eigenfrequencies. In addition we have verified that the codes yield zero frequency modes not only at the maxima but also at the minima of the mass curves.  We give corrected values for the equations of state used by GL, and we add some new equations of state. It is not clear to us what went wrong in their calculations, since for certain EOS our values agree with theirs (EOS C, E, O), for others they differ only slightly (EOS F, L, N), but for some EOS the discrepancy is quite large (EOS A, B, D, G, I).  Additionally we include six more recent equations of state: Two models of \\cite{G85}, one of the model of \\cite{WFF88}, the EOS MPA of \\cite{Wu91}, and two EOS of \\cite{APR98}. Finally, we include three more tables for polytropic equations of state with different polytropic indices. The form we use is given by \\begin{equation} p = \\kappa \\rho^{1 + 1/n}\\:. \\end{equation} In particular, we present for the following values of $\\kappa$ and $n$: ($n = 1, \\kappa = 100\\,$km$^2$), $(n = 0.8, \\kappa = 700\\,$km$^{2.5}$),$(n = 0.5, \\kappa = 2\\cdot 10^5\\,$km$^4)$.  Another interesting feature is the occurrence of avoided mode crossings for realistic EOS. This phenomenon has been thoroughly studied by \\cite{GZ99} for a realistic nucleon EOS and an EOS representing a strange star model. We find that it occurs for all considered realistic EOS, for some it is quite strongly pronounced, for others it is less obvious. ",
        "conclusions": "We have presented updated results for radial oscillations of neutron stars using a quite exhaustive list of currently available equations of state, including some very recent ones. For most equations of state we significantly disagree with the values given by \\cite{GL83}. We have obtained our results by means of two different numerical methods, which agree up to arbitrary precision.  Furthermore, we have checked that our numerical codes yield a zero frequency mode located exactly at both instability points, which are characterized by the local extrema in the mass-density curve. Here, we also obtain full agreement. The overall accuracy, however, is limited by the number of tabulated points for a given equation of state. Here, different numerical interpolation schemes may yield variations in the frequencies up to about three per cent.  Our results agree with the previous results of \\cite{VC92}. However, we use a more complete table for the EOS N (\\cite{S79}), which significantly alters the values one obtains when the table provided by \\cite{LD83} is used.  We have repeated the calculations for the equations of state already used by GL, and we have corrected their given values. In addition, we have included a large number of more recent equations of state. Since most of the present non-linear evolution codes use polytropic equations of state, we also have tabulated the mode frequencies for three different values of the polytropic index $n$."
    },
    "0011/astro-ph0011537_arXiv.txt": {
        "abstract": "We present visibility measurements of the nearby Mira-like star R~Doradus taken over a wide range of wavelengths (650--990\\,nm).  The observations were made using MAPPIT (Masked APerture-Plane Interference Telescope), an interferometer operating at the 3.9-m Anglo-Australian Telescope.  We used a slit to mask the telescope aperture and prism to disperse the interference pattern in wavelength.  We observed in R~Dor strong decreases in visibility within the TiO absorption bands.  The results are in general agreement with theory but differ in detail, suggesting that further work is needed to refine the theoretical models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011471_arXiv.txt": {
        "abstract": "s{ We present a preliminary analysis of the BOOMERanG LDB maps, focused on foregrounds. BOOMERanG detects dust emission at moderately low galactic latitudes ($b > -20^o$) in bands centered at 90, 150, 240, 410 GHz. At higher Galactic latitudes, we use the BOOMERanG data to set conservative upper limits on  the level of contamination at 90 and 150 GHz. We find that the mean square signal correlated with the IRAS/DIRBE dust template is less than 3$\\%$ of the mean square signal due to CMB anisotropy.} ",
        "introduction": "Millimeter-wave emission of interstellar dust can be an important contaminant of sensitive CMB anisotropy surveys (see e.g. \\cite{ang}). However, experimental information on this subject is still coming from serendipitous detections of dust in sensitive anisotropy experiments \\cite{max1} \\cite{mas1}  \\cite{mas2}  \\cite{Leitch} \\cite{Muk} , from the low-resolution surveys of COBE-DMR \\cite{cobe1} and COBE-FIRAS \\cite{cobe3}, and from extrapolation of the full sky surveys of IRAS and DIRBE \\cite{Sch} \\cite{Fink}.  BOOMERanG \\cite{Lan} \\cite{mas3} \\cite{pdb1} is a sensitive millimeter wave telescope operating on a long duration balloon platform. The instrument is designed to take full advantage of long duration stratospheric flights, and features a 1.3 m off/axis, low background, low sidelobes telescope, and ultra sensitive bolometers covering 4 frequency bands (90, 150, 240, 410 GHz) with resolution ranging from 10 to 18 arcmin FWHM.  In the 1998/99 long duration flight, BOOMERanG observed about 2000 square degrees of the sky. About 1300 square degrees are at high galactic latitudes ($b < -20^o$), in a region of sky with the lowest amount of dust emission (constellations of Caelum, Doradus, Pictor, Columba, Puppis). The fluctuation of the 100 $\\mu$m brightness mapped by IRAS is well below 1 MJy/sr in over 500 square degrees, as estimated from the DIRBE-recalibrated IRAS maps \\cite{Sch}.  In this paper we show that, in this region mapped by BOOMERanG, dust contamination in the CMB bands (90 and 150 GHz) is negligible with respect to the CMB signal. The 240 GHz band is partially contaminated by dust emission only in the lowest Galactic latitudes. The 410 GHz map is well correlated with the DIRBE-recalibrated IRAS maps. ",
        "conclusions": "These preliminary results confirm that dust is not expected to be a serious contaminant of sensitive CMB searches, if the region of interest is a low IRAS brightness one, and at frequencies lower than 170 GHz. The quality of the data allows us to extract much more. We are currently working on the determination of dust parameters and on the estimate of the angular power spectrum of fluctuations for this foreground."
    },
    "0011/hep-ph0011308_arXiv.txt": {
        "abstract": "The equation of state for a superconducting cosmic string whose current is due to fermionic zero modes is derived analytically in the case where the back-reaction of the fermions to the background is neglected. It is first shown that the zero mode fermions follow a zero temperature distribution because of their interactions (or lack thereof) with the string-forming Higgs and gauge fields. It is then found that the energy per unit length $U$ and the tension $T$ are related to the background string mass $m$ through the simple relation $U+T= 2m^2$.  Cosmological consequences are briefly discussed. ",
        "introduction": "Topological defects~\\cite{NATO} have been considered in various physical situations, e.g., in the context of condensed matter and cosmology~\\cite{kibble,book}. In many cases of interest in cosmology~\\cite{simulations,BPRS}, they can be approximated as structureless, the relevant dynamics being often assumed not to depend on any specific choice of their internal content. For cosmic defects, this internal content would correspond to the particles that couple to the string-forming Higgs field~\\cite{witten}. However, in the latter example of cosmic strings, it was shown that such a structure might lead to drastic modifications not only of these object dynamics~\\cite{formal,neutral}, which could be seen as a mere academic situation given our present ignorance on their very existence, but also, because of the appearance of new accessible equilibrium states, of the cosmological setting, leading in some instances to actual catastrophes~\\cite{vortons,rdp}. To make a long story short, let us just say that currents imply a breakdown of the Lorentz-boost invariance along the string worldsheet, thereby allowing loop configurations to rotate, the centrifugal force hereby induced having the ability to sustain the loop tendency to shrink because of the tension. The resulting states, called vortons, might be stable even over cosmological timescales, scaling as matter and thus rapidly coming to dominate the Universe evolution, in contradiction with the observations. This leads to constraints on the particle physics theories that predict them at energy scales that are believed to be unreachable experimentally (in accelerators say) in the foreseeable future.  Unfortunately, it appears that the string structure, contrary to their counterparts as fundamental objects~\\cite{superstring}, is not determined by any consistency relation, and is therefore somehow arbitrary, at least at the effective description level~\\cite{witten}. This means in practice that in order to be able to tell anything relevant to (cosmic) string cosmology, one needs to set up a complete underlying microscopic model, arising say, from ones favorite Grand Unified Theory (GUT)~\\cite{davis} or some low-energy approximation of some superstring-inspired model~\\cite{dbp}.  Some generic constructions can however be arranged, as was shown to be the case whence a bosonic condensate gets frozen in the string core~\\cite{formal,neutral}. In such a situation, the boson field phase $\\varphi$, thanks to a random Kibble-like mechanism, may wind along the string itself, thereby producing a current that turns out to be essentially a function of a single state parameter $w$, thus expressible as a phase gradient as \\begin{equation} w \\equiv \\kappa_0 \\gamma^{ab} \\partial_a \\varphi \\partial _b \\varphi,\\label{w}\\end{equation} with indices $a,b,...$ varying within the string worldsheet coordinates defined by the relations \\begin{equation} x^\\mu = X^\\mu_{_{S}} (\\xi^a), \\ \\ \\ \\ \\ \\ \\xi^a \\in \\{\\tau,\\sigma\\},\\end{equation} and $\\gamma^{ab}$ the inverse of the induced metric defined with the background metric $g_{\\mu\\nu}$ as \\begin{equation}\\gamma_{ab} = g_{\\mu\\nu} {\\partial X^\\mu_{_{S}} \\over \\partial \\xi^a} {\\partial X^\\nu_{_{S}} \\over \\partial \\xi^b}.\\label{gammaab}\\end{equation} A straightforward generalization of the Nambu-Goto action is then provided by the $w-$weighted measure as \\begin{equation} {\\cal S} = -m^2 \\int \\hbox{d}^2\\xi \\sqrt{-\\gamma} {\\cal L} \\{w\\},\\end{equation} with $m$ the typical mass scale of symmetry breaking leading to string formation and $\\gamma$ the determinant of the induced metric~(\\ref{gammaab}).  Reasonable microscopic models~\\cite{carpet1} then yield approximate forms for the Lagrangian function ${\\cal L}\\{w\\}$, out of which the dynamical properties of the corresponding strings can be derived~\\cite{larsen,cpg,gpb}.  Such a description remains however essentially classical even though an alternative formalism, also proposed by Carter~\\cite{sigma}, in terms of a dilatonic model, appears more suitable for quantization. This last formalism however, being fully two-dimensional, cannot be used to derive interesting quantities such as the relevant cross-sections for trapped excitations to leave the string worldsheet. This is unfortunate since this is precisely the information one would need for cosmological applications~\\cite{vortons,davis}.  It would therefore seem that by considering fermionic current carriers instead of bosonic ones, one would, because of the intrinsically quantum nature of fermions, obtain a more appropriate description~\\cite{davis,chris}. Besides, fermions are trapped in topological defects because of Yukawa couplings with the string forming Higgs field in the form of zero modes~\\cite{Jackiw}, so that their dynamics is described by simple (although coupled) Dirac equations, which are linear. In the bosonic case, the non-linear (quartic) term is essential in order to ensure the dynamical stability of the condensate so that a solitonic treatment~\\cite{soliton,hp} seems the only way to deal with the underlying quantum physics. This fact dramatically complicates matters and as a result, a complete description yet fails to exist.  However, the fermionic case is not that simple either as here, one faces another technical difficulty for the classical description: it can be shown that there doesn't exist a simple state parameter~\\cite{chris}. An arbitrary spacelike or timelike current can only be built out of at least two opposite chirality spinor fields and will be given by the knowledge of four occupation numbers per unit length. As this is true in particular at least in the zero-temperature limit, we shall be concerned here first with this limit whose validity was assumed to depend on the particular model under consideration~\\cite{davis}. In the following section, we show that setting the temperature to zero is always a good approximation because of the couplings between the fermionic fields and the string-forming Higgs and gauge fields\\footnote{Otherwise, the temperature itself could be taken as a state parameter, so that the usual formalism would be applicable~\\cite{cartertemp}}. These results would seem to imply that the previously derived macroscopic formalism is irrelevant to the fermionic situation (see Ref.~\\cite{many} for a many-parameter formalism). In practice however, as a simple relationship between the energy per unit length and the tension can be found for fermionic currents, the single state parameter formalism can be used that permits to draw some cosmological consequences. ",
        "conclusions": "Fermionic zero modes trapped in cosmic strings are shown to follow a vanishing temperature Fermi-Dirac distribution. This is so because the chirality of the zero modes involved are such that the only possible interaction of the fermions is through gauge boson radiation, leading to an effective loss of energy (on average). As a result, as strings are formed and fermions get condensed along them in the form of zero modes, populating arbitrary high energy levels, they have the possibility to decay radiatively until they reach a zero temperature distribution. Then, as all other interaction terms are identically vanishing, they remain in this state which thus happens to be stable.  Assuming therefore such a vanishing temperature fermionic current-carrying cosmic string, it turns out that the equation of state relating the energy per unit length $U$ and the tension $T$ is of the self-dual~\\cite{carpet1} fixed trace kind, namely $U+T=2m^2$, with $m$ the characteristic string-forming Higgs mass. Although fermionic carriers imply the need of more than one state parameter, this implies that the simplest linear Lagrangian~(\\ref{linear}) provides a good approximation for a classical description of such a vortex. This could in fact have been anticipated as this is the only available equation of state that does not involve any new dimensionfull constant.  Vortons formed with such currents are completely stable, at least at the classical level (see however Ref.~\\cite{davis} for quantum excitations). Assuming backreaction and electromagnetic corrections to be small, one finds that the vorton excess problem~\\cite{vortons} is therefore seriously enhanced for fermionic current-carrier cosmic strings."
    },
    "0011/astro-ph0011317_arXiv.txt": {
        "abstract": "A new method for analyzing point patterns produced by the evolution of gravitational clustering is presented. The method is taken from the study of molecular liquids, where it has been introduced for making a statistical description of anisotropic distributions. The statistical approach is based on the spherical harmonic expansion of angular correlations. A general introduction to the method is given; a theoretical analysis shows that it is partially connected with previous harmonic analyses applied to galaxy catalogs. The effectiveness of the statistical analysis in quantifying clustering morphology is illustrated by applying the statistical estimators to point distributions produced by an ensemble of cosmological $N-$ body simulations with a CDM spectrum. The results demonstrate that the statistical method is able to detect anisotropies in non-random patterns, with different scales being probed according to the expansion coefficients. ",
        "introduction": "\\subsection{ Spherical harmonic analysis} For describing galaxy clustering evolution in the non-linear regime, use is made of different statistical methods, based on the galaxy positions, which have been outlined in the Introduction. The approach introduced here makes use of the positions  of the point distribution within a certain distance from a randomly chosen one.  The method is drawn from molecular dynamics simulations,  where it has been introduced for studying orientational order of supercooled liquids and metallic glasses. A general review can be found in Haile \\& Gray (1980) and Mc Donald (1986), here I will follow the notation of Wang \\& Stroud (1991).  Let us consider a system of $N_p$ particles. The $i-th$ particle has coordinates $\\vec r_i$, in an arbitrary reference frame. For a specified cutoff radius $R_c$ all the particles such that $|\\vec r_i-\\vec r_j|< R_c$ are neighbors of $i$.  The line joining $i$ to one of the $j$ is termed a {\\it bond}.  The angular coordinates of the vector $\\vec \\Delta_{ji}\\equiv \\vec r_j - \\vec r_i $ are $\\theta_j, \\phi_j$ and the quantity \\smallskip \\beq Q_{lm}(\\vec r_i)=\\sum_{j \\ne i} Y_{lm}(\\theta_j,\\phi_j), \\l{eq:defQ} \\eeq is the coefficient of the spherical harmonic expansion of the angular density of the bonds associated with the particle $i$.  In Eq. \\ref{eq:defQ}, and hereafter, summation is understood over all particles $j$ of the distribution such that $|\\vec r_i-\\vec r_j|< R_c$.  The coefficients $Q_{lm}(\\vec r_i)$ are defined as the bond-orientational order parameters (\\cite{wa91}) and they can be drastically changed by a rotation of the reference systems. A natural quantity to consider , which is rotation invariant, is \\beq Q_l(\\vec r_i)=\\sqrt {{{4\\pi}\\over{2l+1}} \\sum_{m=-l}^{m=l} Q^{\\star}_{lm}(\\vec r_i) Q_{lm}(\\vec r_i)}. \\l{eq:Qla} \\eeq  Using the addition theorem for the spherical harmonics, the expression simplifies to  \\beq Q_l(\\vec r_i)=\\sqrt { \\sum_j \\sum_k P_l(\\gamma_{jk})}, \\l{eq:Qlb} \\eeq where $P_l$ is the Legendre polynomial, $$ \\gamma_{jk}\\equiv cos(\\theta_{jk})=\\vec \\Delta_{ji}\\cdot\\vec\\Delta_{ki}/ (|\\vec \\Delta_{ji}||\\vec \\Delta_{ki}|)\\nonumber $$ is the angle between two bonds and  the summations in  Eq. \\ref{eq:Qlb} are then independent of the chosen frame.  The expression  for $Q_l(\\vec r_i)$ can be further simplified if one considers that $\\gamma_{jj}=1$, furthermore it is more convenient to redefine $Q_l(\\vec r_i)$ dividing by the mean number of neighbors $<N_i>=\\bar N_i$ of the particle $i$ :  \\beq Q_l(\\vec r_i)={{1} \\over {\\bar N_i}} \\sqrt { N_i +2 \\sum_j \\sum_{k>j} P_l(\\gamma_{jk})}. \\l{eq:Qlc} \\eeq  For a random distribution the summation terms in the square root are negligible for $N_i \\gg 1$, thus $Q_l(\\vec r_i)\\propto 1 / \\sqrt {N_i}$ for Poisson noise.  For the whole system an order parameter $<Q_l>$ can be defined by averaging over all of the $N_p$ particles: \\beq <Q_l>=\\frac {1}{N_p} \\sum_{i=1}^{N_p} Q_l(\\vec r_i) \\equiv Q_l. \\l{eq:Qld} \\eeq  Thus the clustering distribution can be studied by evaluating the $Q_l$ at a specified cutoff radius. Going to higher coefficients ($ l \\gg 1$) means probing smaller scales in the bond distribution.    However the $Q_l$ do not exhibit spatial information, and a more useful quantity is the auto-correlation $ G_l(r)$  of the coefficients $Q_l(\\vec r_i)$. The function $G_l$ is defined as follows: for all of the $M_p$ pairs $(i,k)$, such that $|\\vec r_i-\\vec r_k|=r \\pm \\Delta r$, where $\\Delta r$ is the thickness of the radial bin, then $G_l(r) $ is the sum over all of these pairs  \\beq G_l(r)={{1}\\over{M_p}} \\sum_i \\sum_k{{4\\pi}\\over{2l+1}} \\sum_{m=-l}^{m=l} Q^{\\star}_{lm}(\\vec r_i) Q_{lm}(\\vec r_i+\\vec r). \\l{eq:GL1} \\eeq  This equation can be greatly simplified: let ${j}$ be the set of neighbors of the particle $i$ and  ${p}$ that of the particle $k$, which satisfy $|\\vec r_j-\\vec r_i|< R_c$  and $|\\vec r_p-\\vec r_k|< R_c$. Then the summation becomes \\begin{eqnarray*} \\sum_{m=-l}^{m=l} Q^{\\star}_{lm}(\\vec r_i) Q_{lm}(\\vec r_i+\\vec r) = \\sum_{m=-l}^{m=l} Q^{\\star}_{lm}(\\vec r_i) Q_{lm}(\\vec r_k)= \\\\ \\sum_{m=-l}^{m=l} \\sum_j Y^{\\star}_{lm}(\\theta_j,\\phi_j) \\sum_p Y_{lm}(\\theta_p,\\phi_p). \\end{eqnarray*} \\smallskip The final expression for $G_l$ is: \\beq G_l(r)={{1}\\over{M_p}} \\sum_i \\sum_k \\sum_j \\sum_p P_l(\\Gamma_{jp}), \\l{eq:GL4} \\eeq with {\\bf $\\Gamma_{jp}$} being the angle between $\\vec \\Delta_{ji}$ and $\\vec \\Delta_{pk}$.  The summation over the pairs  is $\\sum_i \\sum_k$, with the sum over the particles $k$ only for those particles with $|\\vec r_i-\\vec r_k|=r \\pm \\Delta r$.  Note that now it is not possible to simplify the equation as was done for Eq. \\ref{eq:Qlc}, because one is considering the correlation between {\\it different} order parameters.  The information about the degree of angular correlation between bonds contained in the $G_l(r)$ is more explicit if the $G_l$ themselves are divided by  the zero-order correlation function \\beq G_0(r)={{4\\pi}\\over{M_p}} \\sum_i \\sum_k Q^{\\star}_{0}(\\vec r_i) Q_{0}(\\vec r_i+\\vec r). \\l{eq:GL2} \\eeq The functions $G_l(r)$ can be made dependent also on the direction of $\\vec r$ by considering all the pairs with orientation of the relative separation $\\vec r$ in the solid angle range $\\Omega_{\\vec r} \\pm \\Delta \\Omega_{\\vec r}$. These functions $G_l(\\vec r)$ are of scarce practical use because of the noise level which is introduced when applied to a finite sample.    \\subsection { Theory} The equations that have been obtained for $Q_l$ and $G_l(r)$ are defined through a point set distribution; from  a theoretical point of view it is possible to relate these functions to the power spectrum of the density fluctuations. Spherical harmonic analysis has been widely applied in cosmology, both to the galaxy distribution (\\cite{sc92}; \\cite{sc93}; \\cite{ba95}) and to the study of cosmic background radiation maps (\\cite{hu97}; \\cite{ba99}). I will follow here the notation of Scharf et al. (1992), who have analyzed the angular distribution of the IRAS redshift survey using spherical harmonics. For the generic distribution considered in Sect.2.1 let us define  $n(\\vec r)$ to be the number density and $\\bar n$ to be the average one. It is also useful to introduce $\\Phi(r)$ as a generic selection function, here $\\Phi=\\bar n$.  Then the coefficients $Q_l^m(\\vec r_o)$, $\\vec r_o$ being the center position, are given by $$ Q_l^m(\\vec r_o)=\\int \\Phi(\\vec r^{\\prime}) {{n(\\vec r^{'})} \\over {\\bar n}} Y_l^m(\\Omega_{\\vec r^{'}}) d^3 r^{'} ~~~~\\vec r=\\vec r_o+\\vec r^{\\prime}, $$ where the integral is taken over a sphere of radius $R_c$. The number density is $$ n(\\vec r)=\\bar n [ 1+ \\delta(\\vec r)]=\\bar n[1+\\delta(\\vec r_o+\\vec r^{'})], $$ here $\\delta(\\vec r)$ is the density fluctuation. Thus the equation for $Q_l^m(\\vec r_o)$ becomes \\bea Q_l^m(\\vec r_o)& =& \\int \\Phi [1+\\delta] Y_l^m d^3 r^{'}= \\nonumber \\\\ & & \\int \\Phi Y_l^m d^3 r^{\\prime}+\\int \\Phi(\\vec r^{'}) \\delta (\\vec r_o+\\vec r^{'}) Y_l^m d^3 r^{'}, \\nonumber \\eea where the first term on the rhs is zero for $l \\ge 1$. The density fluctuation $\\delta(\\vec r)$ can be expanded into Fourier modes, and the final expression for $Q_l^m(\\vec r_o)$  is  \\beq Q_l^m(\\vec r_o)=\\sum_{\\vec k} \\delta_{\\vec k} e ^{ i\\vec k \\vec r_o} \\int \\Phi(\\vec r^{'}) e^{i \\vec k\\cdot r^{'}}  Y_l^m(\\Omega_{\\vec r^{'}}) d^3 r^{'}. \\eeq The $ e^{i\\vec k \\vec r }$ is now decomposed into spherical functions, so that the ensemble average gives \\begin{eqnarray} < |Q_l^m|^2 >& = &( 4\\pi)^2 \\sum_{\\vec k} |\\delta_{k}|^2 \\left( \\int_0^{R_c} \\Phi(x) x^2 j_l(kx) dx \\right) ^2 \\nonumber \\\\ & & |Y_l^m(\\Omega_{\\vec k})|^2, \\l{eq:qlm} \\end{eqnarray} here the $j_l(x)$s are the spherical Bessel functions. The summation over $k$ can be translated into an integral  and $\\int |Y_l^m|^2 d\\Omega=1$, so that \\bea < |Q_l^m|^2 >& =& V {{2} \\over {\\pi}} \\int_0^{\\infty} P(k)k^2 dk \\left[ \\int_0^{R_c} \\Phi(x) x^2 j_l(kx) dx \\right] ^2 \\nonumber \\\\ & & + {\\cal N}, \\l{eq:QF} \\eea where $P(k)=<|\\delta_k|^2>$, $V$ is a normalizing volume and $\\cal N$ is the shot noise term ${\\cal N} =\\int \\Phi(x)x^2 dx$.  The coefficients $C_l$ are defined as \\beq C_l^2={ {1} \\over { {\\cal N} (\\Omega/4\\pi)(2l+1)}} \\sum_m |Q_l^m|^2, \\l{eq:Cla} \\eeq and assuming statistical isotropy they can be related to Eq.(\\ref{eq:Qld}) by $Q_l\\simeq C_l/ {\\bar N_i^{1/2}}$. For an uncorrelated distribution $C_l=1$, furthermore the integral $ \\int_0^{R_c} \\Phi(x) x^2 j_l(kx) dx \\equiv\\Psi_l(kR_c)$ defines a bidimensional window function in the parameter space $(R_c,l)$. Increasing the value of $l$ means probing smaller angular scales, the value of  the cutoff radius $R_c$ sets the size of the spectrum contribution to the integral (see,  for example, Fig. 4 of \\cite{sc93}). The statistical method introduced in Sect.IIA, and originally proposed in a different field, is then mathematically equivalent to the spherical harmonic analysis which has been applied to galaxy surveys (\\cite{sc93}). The order parameter $Q_l$ is just related to the angular power spectrum estimator $C_l$ (\\cite{sc92}).  However Eq. (\\ref{eq:qlm}) is valid as long as the position of the observer located at $\\vec r_0$ can be considered as random with respect to the bond distribution inside the sphere of radius $R_c$. This is a good approximation for values of the cut-off radius  $R_c $ much larger than the 2-point clustering length, a condition which is valid for the analyzed angular catalogs.  In Sect. 4 the statistical method is tested using the point distributions obtained from a set of cosmological simulations, and the coefficients $Q_l$ are calculated for different parameters and cut-off radii $R_c$. Therefore one has to consider in principle also the three-point correlation terms  between the particle located at $\\vec r_0$ and the \\\\ other particles in the bond sphere. The ensemble average for $ |Q_l^m|^2 $ is then (Peebles 1980, sect. 36B):  \\bea < |Q_l^m|^2 > & = &\\int \\Phi d \\vec x  \\int \\Phi d \\vec y [ 1+\\xi(x) + \\xi(y) +\\xi(|\\vec x - \\vec y|) \\nonumber \\\\ &  & + \\zeta(x,y,|\\vec x - \\vec y|) ] Y_l^m(\\Omega_{\\vec x}) {Y^{\\star}_l}^m(\\Omega_{\\vec y}), \\l{eq:Qg3} \\eea where $\\vec x$ and $\\vec y$ are the particle separations relative to the origin and $\\zeta$ is the reduced three-point correlation function. Integration over solid angles removes for $l\\geq1$ the first three terms on the rhs of the equation. Of the remaining terms, the first ( $ \\xi(|\\vec x - \\vec y|= \\sum_{\\vec k}  |\\delta_k|^2 e^{i \\vec k( \\vec x - \\vec y)})$, gives Eq. (\\ref{eq:qlm}), while the other takes into account for the triangle configuration the  correlations between the three particles. The integral can be evaluated in terms of the Fourier transform of $\\zeta$ (\\cite{fr84}), the bispectrum $B(\\vec k_1, \\vec k_2, \\vec k_3)$ :  \\bea \\int \\Phi d \\vec x  \\int \\Phi d \\vec y \\zeta(x,y,|\\vec x - \\vec y|)Y_l^m(\\Omega_{\\vec x}) {Y^{\\star}_l}^m(\\Omega_{\\vec y}) =  & & \\nonumber \\\\ \\sum_{\\vec k_1 \\vec k_2 \\vec k_3} B(\\vec k_1, \\vec k_2, \\vec k_3) \\int \\Phi d \\vec x  \\int \\Phi d \\vec y Y_l^m(\\Omega_{\\vec x}) {Y^{\\star}_l}^m(\\Omega_{\\vec y}) & & \\nonumber \\\\ e^{i(\\vec k_1 \\vec x_1+ \\vec k_2 \\vec x_2 + \\vec k_3 \\vec x_3)}= \\sum_{ \\vec k_2 \\vec k_3} B(\\vec k_1=-(\\vec k_2+\\vec k_3), \\vec k_2, \\vec k_3) &  &\\nonumber \\\\ \\int \\Phi d \\vec x  \\int \\Phi d \\vec y Y_l^m(\\Omega_{\\vec x}) {Y^{\\star}_l}^m(\\Omega_{\\vec y}) e^{i( \\vec k_2 \\vec x + \\vec k_3 \\vec y)} =  & & \\nonumber \\\\ (4 \\pi )^2  (-)^l \\sum_{ \\vec k_2 \\vec k_3} \\int \\int B(\\vec k_1, \\vec k_2, \\vec k_3) \\Psi_l(k_1R_c)   \\Psi_l(k_2R_c) & & \\nonumber \\\\ Y_l^m(\\Omega_{\\vec k_1}) {Y^{\\star}_l}^m(\\Omega_{\\vec k_2})~, ~~~~~~~~~~~~~~~~~~~~~~~~~& & \\l{eq:bs} \\eea where the constraint $\\sum \\vec k_i=0$ is required by homogeneity, $\\vec x \\equiv \\vec x_2 - \\vec x_1$ and $\\vec y\\equiv \\vec x_3 -\\vec x_1$. In second order perturbation theory, for Gaussian initial conditions, the bispectrum can be expressed as ( Fry 1984 ; \\cite{ve20}) : \\beq B(\\vec k_1, \\vec k_2, \\vec k_3)= K( \\vec k_1, \\vec k_2) P(k1) P(k2) + cyc.~, \\l{eq:bs1} \\eeq here $cyc.$ means cyclic terms. The term $K(\\vec k_1, \\vec k_2)$ is weakly dependent on the cosmology and can be written as (Fry 1984; Goroff et al. 1986; \\cite{ve20}): \\beq K( \\vec k_1, \\vec k_2)= A_0+A_1 cos(\\theta_{12})+A_2 cos^2(\\theta_{12}), \\l{eq:bs2} \\eeq where the coefficients $A_i$ depend on the assumed biasing expansion for the local density field and $\\theta_{12}$ is the angle between $\\vec k_1$ and $\\vec k_2$. The integral (\\ref{eq:bs}) can then be evaluated for a particular cosmological model. The dependence of $B$ on the triangle shape in the $k$-space shows that the integral is non zero only for $l \\leq2$. The integral (\\ref{eq:bs}) is a particular case of the integrals considered by Verde, Heavens and Matarrese (2000). These authors have expanded the projected galaxy density in spherical harmonics and studied the 3-point function of the expansion coefficients, which is a quantity directly related to the bispectrum. The evaluation of the above integral is rather complicated (see, e.g., Eq. (28) of Verde, Heavens and Matarrese 2000) and will not be considered here. In Sect. 4 it will be seen that the corrections (\\ref{eq:bs}) to Eq. (\\ref{eq:qlm}) must be rather small already for $R_c \\simeq 2 r_0$.   This theoretical analysis can also be used to derive an expression for the quantity $G_l(\\vec r)$, which contains spatial information about the bond angular distribution. The evaluation of $G_l(\\vec r)$ now involves correlation terms up to the fourth order. In analogy with Eq. (\\ref{eq:Qg3}) one has to compute for  $G_l(\\vec r)$ the ensemble average \\bea \\lefteqn{ G_l(\\vec r)=<Q_l^m(\\vec r_o) Q_l^{\\star m}(\\vec r_o+\\vec r) > = } \\nonumber \\\\ \\,& \\int \\Phi d \\vec x  \\int \\Phi d \\vec y [ \\xi(|\\vec r+ \\vec y - \\vec x|)+ \\zeta(x,|\\vec r + \\vec y - \\vec x |, |\\vec r + \\vec y|) \\nonumber \\\\ \\,&+\\zeta(|\\vec r + \\vec y - \\vec x |, y, |\\vec r - \\vec x|) + \\xi(|\\vec r + \\vec y|) \\xi(|\\vec x -\\vec r|)+ \\xi(r)\\nonumber \\\\ \\,&\\xi(|\\vec r + \\vec y - \\vec x)|)+ \\eta (x,r,y,|\\vec r + \\vec y|,|\\vec r - \\vec x|, |\\vec r + \\vec y - \\vec x|)]\\nonumber \\\\ \\,&{ Y_l^m(\\Omega_{\\vec x}) {Y^{\\star}_l}^m(\\Omega_{\\vec y}),~~~~~~~~~~~~ ~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~} \\l{eq:Ga1} \\eea where $\\eta$ is the reduced four-point correlation function, and only the terms which are non-zero after the integration over the solid angles are shown in the integrals; the separations of the four-particle configuration of Eq. (\\ref{eq:GL4}) are defined as:~$ \\vec r= \\vec r_k - \\vec r_i  ,~ \\vec y= \\vec r_p - \\vec r_k , ~ \\vec x = \\vec r_j - \\vec r_i$. The integrals in (\\ref{eq:Ga1}) can be evaluated using the same method adopted  for the integral (\\ref{eq:bs}). As previously discussed, from the results of Sect. 4 for $< |Q_l^m|^2 >$ , the third and fourth-order terms are expected to be negligible. I will consider here only the leading order term $ \\xi(|\\vec r+ \\vec y - \\vec x|)$ and the ensemble average for $G_l(\\vec r)$ is then  \\bea \\lefteqn{ <Q_l^m(\\vec r_o) Q_l^{\\star m}(\\vec r_o+\\vec r) > = } \\nonumber \\\\ \\,& \\sum_{\\vec k} e^{-i \\vec k \\vec r} |Y_l^m(\\Omega_k)|^2 \\left [ \\int \\Phi j_l(kx)x^2 dx \\right]^2 (4\\pi)^2 |\\delta_k|^2. \\l{eq:Ga} \\eea  The summation over $\\vec k$ becomes an integral and if $e^{-i \\vec k \\vec r}$ is now decomposed as \\beq e^{-i \\vec k \\vec r} = 4 \\pi \\sum_{pq} i^p j_p(kr) Y^*_{pq}(\\Omega_{\\vec r})Y_{pq}(\\Omega_{\\vec k}), \\eeq then the triple integral for the spherical harmonics $Y_l^m$ in Eq.(\\ref{eq:Ga}) is \\begin{eqnarray*} \\lefteqn {\\int Y_{pq} Y_{lm}^* Y_{lm} d\\Omega_k=\\int Y_{pq}(-)^mY_{l-m}Y_{lm}d\\Omega_k= } \\nonumber \\\\ \\,& (-)^m \\left [ {(2l+1)}^2~({2p+1)} \\over {4\\pi} \\right ]^{1/2} \\left ( \\begin{array}{lcr} l & l & p \\\\ 0 & 0 & 0 \\end{array} \\right ) \\left ( \\begin{array}{lcr} l & l & p \\\\ m & -m & q \\end{array} \\right ). \\end{eqnarray*} The symbols $ \\left ( \\begin{array}{lcr} l_1 & l_2 & l_3 \\\\ m_1 & m_2 & m_3 \\end{array} \\right ) $ are the Wigner $3-j$ coefficients. Because of their properties, the summation over the $q$ terms in Eq. (\\ref{eq:Ga}) is non-zero only for $q=m-m=0$, furthermore the $p$ summation admits only terms with $2l+p=even, 0 \\leq p \\leq 2l$. Finally the expression for $G_l(\\vec r)$ becomes : \\bea G_l(\\vec r)  =  { {4\\pi}\\over {2l+1}} \\sum_m <Q_l^m(\\vec r_o) Q_l^{\\star m}(\\vec r_o+\\vec r) >=  \\nonumber  \\\\ {{(4\\pi)^4}\\over {(2\\pi)^3}} V \\int_0^{\\infty} dk k^2 \\left [ \\int_0^{R_c} \\Phi x^2 j_l(kx) dx \\right]^2 P(k) \\nonumber  \\\\ \\sum_{p( even )=0}^{2l} i^p j_p(kr) Y_{p0}^{\\star}(\\Omega_{\\vec r}) \\sum_{m=-l}^{m=l} (-)^m {{(2p+1)}^{1/2} \\over {(4\\pi)^{1/2}}} \\nonumber \\\\ \\left ( \\begin{array}{lcr} l & l & p \\\\ 0 & 0 & 0 \\end{array} \\right ) \\left ( \\begin{array}{lcr} l & l & p \\\\ m & -m & 0 \\end{array} \\right ). \\l{eq:Gb} \\eea In the above equation  the function $G_l$ still depends on the orientation of $\\vec r$ through the $m$ terms. However, owing to statistical isotropy, this dependence should not be present. In fact, it can be shown that (\\cite{ga95}) $ \\sum_m (-)^m \\left ( \\begin{array}{lcr} l & l & p \\\\ m & -m & 0 \\end{array} \\right ) =(2l+1)^{(1/2)}(-)^l \\delta_{p,0}, $ and Eq. (\\ref{eq:Gb}) simplifies to  \\bea \\lefteqn { G_l(r) ={ {(4\\pi)}^3 \\over {(2\\pi)^3}} V \\int_0^{\\infty} dk P(k) k^2 } \\nonumber \\\\ \\,& \\left[ \\int_0^{R_c} \\Phi(x) x^2 j_l(kx) dx \\right] ^2 j_0(kr). \\l{eq:Gc} \\eea  From this function one can recover the standard definition for $\\xi(r)$ taking for $l=0$ the limit $ R_c \\rightarrow 0 , \\int_0^{R_c} \\Phi(x) x^2 j_0(kx) dx \\rightarrow 1/4\\pi$. The statistical descriptor of clustering defined by Eqs. (\\ref{eq:Qld}) and (\\ref{eq:GL4}), when applied to a spatial point process, should be able to describe in a quantitative way the patterns generated by the clustering distribution. As outlined in the Introduction different methods have been employed to this end; the most important advantage of the order parameters introduced here is that they can can be used to study spatial patterns in a well-defined way by varying the angular scale $l$ or the depth $R_c$ of the bond spheres. In order to test the clustering features described by the order parameter $Q_l$ and correlation function $G_l$ , I have then analyzed point distributions generated by a set of cosmological $N$-body simulations with CDM spectra. The simulations are evolved in time starting from an early epoch when linear theory applies. The clustering evolution gives rise  to a variety of spatial structures and the statistical descriptors are then applied, in order to analyze the produced patterns, to numerical outputs selected at different epochs. The results obtained are discussed in Section 4. ",
        "conclusions": "The order parameters $Q_l$ are shown in Figs. \\ref{Qa},~\\ref{Qb} \\& \\ref{Qc} for different values of the expansion factor $a(t)$, threshold levels and cut-off radius  $R_c$. In Fig. \\ref{Qa} the time evolution of $Q_l$ is plotted as a function of $l$ for different epochs. Several generic features can be seen in the plot: the coefficients $Q_l$ grow with time as the clustering develops and decay with increasing $l$ as smaller angular scales are probed. This follows directly form the shape of $P(k)$. For a random test distribution  with $N_i \\simeq 10^3$ neighbors $Q_l \\simeq const \\simeq 1 / \\sqrt {N_i}$, and the estimator (\\ref{eq:Qld}) is clearly able to detect non-random distributions with a high level of significance. The scatter in the numerical ensemble  can be judged from the size of the error bars, for the sake of clarity they have been shown in each plot only for a single case. For the plot of Fig. \\ref{Qa} the point distribution is that obtained from the $\\{x\\}_{-1}$ set. In this case the computation of the $Q_l$ can be greatly reduced by using a random sub-sample of the particle set. This has been done also for other sets $\\{x\\}_{s}$, according to required computational load needed to evaluate the coefficients $Q_l$.  \\begin{figure} \\centerline{ \\psfig{file=plqi2.ps,height=4.in,width=3.8in} } \\caption{ Here the estimators $Q_l$ have been computed at a fixed time and cut-off radius $R_c$ from the point distributions generated by $\\{x\\}_{s}$, with $s=-1,0.5,1.3,2.5$. The black squares are the rescaled linear theory coefficients $C_l$, computed from the CDM power spectrum at $a(t)=6$ and for $R_c=20$. The power spectrum has been normalized as in the simulations, so that the points compare directly with the $\\delta=-1$ $Q_l$ coefficients. The black triangles refer to the same theoretical coefficients, but computed in this case using the non-linear power spectrum of the simulation.}  \\label{Qb} \\end{figure}  The results obtained have been found to be quite robust to changes in the size of the random sub-sample, usually an accurate evaluation of $Q_l$ can be obtained already with a $ \\simeq 20 \\% $ random sub-sample of the particle set. How the $Q_l$ change when particle populations with different clustering properties are selected is shown in Fig. \\ref{Qb}. For a fixed $a(t)$ and $R_c$ the coefficients $Q_l$ have been evaluated for particle subsets associated with different thresholds. For any $l$ the coefficients $Q_l$ get larger when higher thresholds $\\nu$ are selected. The choice of higher $\\nu's$ corresponds to a more clustered distribution, and this is clearly detected by the coefficients $Q_l$. It is worth stressing that the trend is obtained according to the adopted definition (\\ref{eq:Qlc}) for $Q_l(\\vec r_i)$. The angular power spectrum coefficients $C_l$ scale approximately as $\\propto N_i$, thus the enhanced clustering of particle subsets associated with higher $\\nu's$ is masked for the $C_l$ by the drop in $N_i$, with respect lower thresholds. The black squares in Fig. \\ref{Qb} are the rescaled linear theory coefficients (\\ref{eq:Cla}). In this case the comparison has been made for the population $\\delta \\geq -1$, in order to avoid errors introduced by a linear bias approximation for Eq. (\\ref{eq:Cla}), if subsets $\\{x\\}_s$ corresponding to higher thresholds $\\nu$ would have been selected.   \\begin{figure} \\centerline{ \\psfig{file=plq3.ps,height=4.in,width=3.8in} } \\caption{ The coefficients $Q_l$ are shown at $a(t)=6$ for four different cut-off radii $R_c=10,20,40,60$ ( in $h^{-1}Mpc$ units). The estimators have been applied, as in Fig. \\ref{Qa}, to the whole particle distribution.} \\label{Qc} \\end{figure}  There is a good agreement with the numerical experiment up to $l \\simeq 5$; beyond this angular scale ( $\\simeq \\pi/5$) non-linearity effects make the $Q_l$ depart from the linearized coefficients. For the adopted normalization the size of the non-linear scale is about $R_{NL} \\simeq 5-8 h^{-1} Mpc$, at the epoch shown in the plot, which is in rough agreement with $R_c \\Delta \\theta$. When the power spectrum of the simulation particles is used to compute the theoretical coefficients there is an excellent agreement with the corresponding measured $Q_l$ down to the smallest angular scale (black triangles of Fig. \\ref{Qb}). For $ l \\leq 2$ the small differences show that the corrections (\\ref{eq:bs}) to Eq. (\\ref{eq:qlm}), for the considered cosmological model are negligible already for $ R_c \\simgt 20h^{-1}Mpc$.  The coefficients $Q_l$ are plotted for different cut-off radius $R_c$, in Fig. \\ref{Qc}, for a fixed time and a given set $\\{x\\}_s$. There is a general tendency for the $Q_l$ to get smaller at a fixed $l$ when $R_c$ is increased. This is expected because the bond distribution approaches  isotropy as $R_c$ gets higher. This trend is confirmed for all values of $l$ and choices of different distributions $\\{x\\}_s$. Note also how the scattering in ensemble is reduced ($R_c=40 h^{-1}Mpc$) with respect to the smaller values of $R_c$ shown in the previous plots. The analysis performed so far shows that the order parameters $Q_l$ can be used as reliable statistical indicators of the global clustering properties of a point distribution.  From the histograms of Fig. \\ref{Qc} an interesting result is that the quadrupole coefficient $Q_2$ of the bond angular distribution is larger than the dipole coefficient $Q_1$. This is valid with a high significance only for values of the cut-off radii $R_c \\simlt 2 r_0$, close the value of the 2-point clustering length. This excess of local anisotropies, compared to what is found for $R_c \\simgt 20 h^{-1} Mpc$, suggests that a significant amount of substructure is present in the particle distribution of dark matter halos. Jing (2000) has analyzed the halo dark matter density profiles from a set of cosmological $N-$body simulations. For those halos which are substructure rich and less virialized, he found large deviations from the analytical profile of  Navarro, Frenk \\& White (1995). The above results therefore suggest that the proposed statistical method can be profitably used to correlate in a quantitative way the local halo anisotropies with other halo properties, when investigating their evolution in different cosmological models.   \\begin{figure} \\centerline{ \\psfig{file=plg1.ps,height=4.in,width=3.8in} } \\caption{ The bond correlation functions $G_l(r)$, divided by $G_0$, are shown at $a(t)=6$ and $R_c=40$ for the distribution obtained from $\\{x\\}_{s=0.5}$. The different histograms are for different values of $l$, only $G_l$ up to $l=4$ are plotted. The line with a value for $G_l$ close to zero is for a random distribution.} \\label{Ga} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=plg11.ps,height=4.in,width=3.8in} } \\caption{ As in Fig. \\ref{Ga}, but in this case $G_l(r)$ have been evaluated summing in (\\ref{eq:GL4}) over all the neighbors of the center particles; in the previous figure $G_l(r)$ were calculated taking random sub-sample of the neighbor sets.} \\label{Ga1} \\end{figure}  To analyze the clustering shape, more interesting results are obtained if the estimator (\\ref{eq:GL4}) is applied to point distributions generated by different $\\xs$.  The total operation count of applying Eq. (\\ref{eq:GL4}) to a given point distribution scales as $\\propto N_p^4$, thus making it prohibitive to evaluate $G_l(r)$ for the whole set $\\{x\\}_s$  of the simulation box. Therefore one has to resort to compute $G_l$ taking as centers $i$ in Eq. (\\ref{eq:GL4}) a random subsample of $\\{x\\}_s$. The summations \\\\over the particles $k$ at a distance $r$ from $i$ and the neighbor particles of the chosen centers ($\\sum_j \\sum_p$) are instead run over the whole set $\\{x\\}_s$, although a random dilution of the neighbor sets was performed in many cases in order to reduce the total computational cost. For a given particle $i$ taken as a center the mean number $\\bar N_i$ of neighbors $j$ within a distance $R_c$ from $i$ ranges from $\\simeq10^2$ to $\\simeq10^4$, according to the chosen cut-off radius $R_c$, threshold level $\\nu$ and expansion factor $a(t)$.   However the use of a random subset requires some care with respect to the previous procedures used to compute the coefficients $Q_l$. It has been found that the function $G_l(r)$ is subject, in some cases, to large sample-to-sample variations when different random subsets are chosen from a given $\\{x\\}_s$ to evaluate the corresponding $G_l$. This shows that $G_l(r)$ is a sensitive measure of the clustering patterns generated by a point distribution.  \\begin{figure} \\centerline{ \\psfig{file=plg2.ps,height=4.in,width=3.8in} } \\caption{ The functions $G_l(r)$ plotted here have been obtained from different point distributions and the same values of $l,R_c$ and $a(t)$ (given in the panel). The distributions are those given by the simulation particle subsets $\\{x\\}_{s}$, with $s=0.5,1.3,2.5$, and  $\\{x\\}_{s=-1}$.} \\label{Gb} \\end{figure}  One of the consequences of this result is that the calculated $G_l(r)$ analyze in this case the clustering morphology of the point distribution as it is measured according to the selected random sub-sample, rather than that of the whole $\\{x\\}_s$. In order to consistently compare $G_l(r)$ for \\\\different choices of several parameters, the same random subset has been chosen at different times and for different \\\\$R_c$. Furthermore $G_l(r)$ for point distributions obtained for different thresholds have been calculated with random subsamples subject to the constraint that for a given $\\xs$ the random subset was part of the random sample corresponding to a lower threshold, and so on down to the lowest level with $s=-1$. The application of these constraints leaves about $8,000$ particles in the simulation volume, which have been taken as center particles when the $G_l(r) $ were calculated. This procedure allows a consistent analysis to be made of the clustering morphologies detected by the $G_l(r)$ at different epochs and different values of the cut-off parameter $R_c$, angular scale $l$ and threshold $\\nu$. The results obtained are shown in the following plots. Hereafter $G_l(r)$ is the estimator (\\ref{eq:GL4}), applied to point distributions, divided by $G_0(r)$. Because of the large number of cases spanned  in the parameter space, only a few of them have been plotted, the main features seen are quite general.   In Fig. \\ref{Ga} $G_l(r)$ is given as a function of $r$ for different values of $l$, up to $l=4$. The values for the other parameters are given in the panel. The correlation between angular distributions for different bonds decreases with increasing $r$ and as smaller scales are probed, but with some exceptions. In this case $G_l$, for $l=1$, has a significant correlation above the noise extending for a wide range of spatial distances. In other cases a similar trend has been found also for $l=2$. $G_l(r)$ is a measure of the degree of coherence in the clustering patterns. For the example shown in the figure the result is a web of large-scale structures, in the examined point distribution, which is detected by the estimator. These features are detected at the lowest order only ($l\\leq 2$), since their distribution in the bond spheres is limited to large angles. For $r \\rightarrow0$ the correlation strength increases because the two neighbor spheres overlap and partially trace the same structures. The continuum limit (\\ref{eq:Gc}) is reached in the large sample regime $\\bar N_i \\rightarrow \\infty$. For a finite number of points $G_l(r)$ measures the anisotropies of the distribution which samples the density field.   A theoretical estimate of how $G_l(r)$ are affected by noise terms is difficult, therefore an indirect check has been performed for the examples showed in the plots evaluating $G_l(r)$ in correspondence of a random distribution. For each of the considered cases this was obtained by considering tha same set of $M_p$ pairs used in Eq. (\\ref{eq:GL4}) to evaluate $G_l(r)$, but with a random distribution of $N_r=10^3$ neighbors within each of the two spheres of radius $R_c$ centered at the pair coordinates. The results show that $G_l \\simeq 0$ for a random distribution of neighbors, so that the measured $G_l(r)$ are well above the noise.  For the example showed in Fig. \\ref{Ga} $G_l(r)$ have been computed using random sub-samples of the $\\bar N_i\\simeq 10^4$ neighbors of the center particles $i$. In this case the fraction $f$ of selected neighbors was $f \\approx 1/4$. Finite sample effects on the measured $G_l(r)$ can be estimated by repeating the evaluation of $G_l(r)$ for the case of Fig. \\ref{Ga}, but this time running the summations in (\\ref{eq:GL4}) over all the neighbors particles. The results are showed in Fig. \\ref{Ga1}, a comparison with the previous figure shows that the $G_l(r)$ profiles are qualitatively similar, suggesting that measurements of the clustering pattern by $G_l(r)$ are not affected by a dilution of the neighbor sets. Empirically it has been found that for the examples showed here this is valid for $ f\\bar N_i \\simgt 10^2$.   $G_l(r)$ is also sensitive to different degrees of clustering, as shown in Fig. \\ref{Gb}. In this case the $G_l$ have been computed for different $\\xs$ with the other parameters being kept constant (see the panel).  \\begin{figure} \\centerline{ \\psfig{file=plg3.ps,height=4.in,width=3.8in} } \\caption{ The $G_l(r)$ are shown for different values of the cut-off radius: $R_c=10,20,30,40h^{-1}Mpc$. The  values of the other parameters are given in the panel.} \\label{Gc} \\end{figure}  It is important to stress that this result was obtained because of the way in which the random subsets were chosen; if different subsets had been used for different $\\xs$, then sample-to-sample variations would have dominated over the differences in $G_l$ for different thresholds. A generic feature of $G_l$ is a decrease in the signal as $R_c$ is increased (Fig. \\ref{Gc}). This is expected, since in going to higher cut-off radii, the point distribution in the bond sphere approaches  isotropy.  As previously stressed $G_l(r)$ has been found to be very sensitive to sample variations, the size of the fluctuations can be checked in Fig. \\ref{Gd}, where $G_l(r)$ for a specified set of parameters has been computed for the different $\\xs$ of the numerical ensemble. Because of the computational cost of  evaluating $G_l(r)$ I have calculated these functions for the whole simulation ensemble only for $a(t)=6$. The plot shows clearly how fluctuations in the ensemble, in this case, are of the same order as the variations in $G_l$ shown in the previous plots when the input parameters were changed.  \\begin{figure} \\centerline{ \\psfig{file=plg4.ps,height=4.in,width=3.8in} } \\caption{ The correlation  functions $G_l(r)$ are plotted here  for $l=1,R_c=40,a(t)=6$ and point distributions obtained from $\\{x\\}_{s=2.5}$. Different histograms refer to $\\{x\\}_{s}$  from simulations with different random realizations of the ensemble.} \\label{Gd} \\end{figure}  In one simulation large-scale correlations are clearly detected for $r\\simgt60 h^{-1} Mpc$. These results are in accordance with those of Doroshkevich, Fong \\& Makarova (1998), who have analyzed the evolution of large-scale structure in standard CDM $N$-body simulations. Their analysis of sheets and filaments in the particle distribution is based upon the 'core sampling' method (Buryak, Doroshkevich \\& Fong 1994). A proper comparison is difficult because of the different statistical methods employed. The results obtained by Doroshkevich, Fong \\& Makarova (1998) reveal a richness of structures on scales of $\\sim10 -100h^{-1} Mpc$, in agreement with the correlation range measured here by the functions $G_l$.   In Fig. \\ref{Ge} the functions $G_l(r)$ are shown, for different realizations, for the same example as in Fig. \\ref{Gd}, but in this case for $\\{x\\}_{s=0.5}$ instead of $\\{x\\}_{s=2.5}$. The scatter in the ensemble is clearly reduced; this is a consequence of the lower threshold considered. For a point distribution approaching that of the whole mass in the simulation volume then $G_l(r)$ has a larger contribution from particles which are not part of filaments or of other large-scale structures. The signal is thus reduced as also are the variations between different random realizations. This is valid for cut-off radii much larger than the non-linear scale. If $R_c=10h^{-1}Mpc$ then the functions $G_l$ would have shown no sign of correlations, with strong fluctuations around zero independently of the separation distance $r$. The scattering in the ensemble is also reduced when smaller angular scales are probed (Fig. \\ref{Gf}).  \\begin{figure} \\centerline{ \\psfig{file=plg5.ps,height=4.in,width=3.8in} } \\caption{ The same as in Fig. \\ref{Gd}, but for distributions obtained from $\\{x\\}_{s=0.5}$.} \\label{Ge} \\end{figure}  To summarize, the observed results have shown that the statistical method defined by the functions $Q_l$ and $G_l(r)$ can be used to analyze the clustering morphology produced by gravitational clustering in a quantitative way. The method was originally devised in a completely different field for studying anisotropic structures and its application to cosmological clustering shows that, to some extent, it overlaps with previous analyses. The usefulness of the method has been tested by applying it to a set of cosmological N-body simulations with a CDM power spectrum. An application of the estimators to selected numerical outputs shows that the statistic is clearly able to discriminate between particle populations with different degrees of clustering. Several shape statistics have been introduced to quantify the presence of filaments or sheet-like structures in the clustering network. The function $G_l(r)$ used here describes anisotropies in the clustering distribution by measuring the degree of correlation between the angular densities as seen from two different observers separated by $r$. $G_l(r)$ can then be considered a statistical measure of clustering patterns, with different scales probed by varying the input parameters $l$ and $R_c$. This provides an alternative to other methods proposed so far to quantify the clustering morphology. \\begin{figure} \\centerline{ \\psfig{file=plg6.ps,height=4.in,width=3.8in} } \\caption{ As in Fig. \\ref{Gd}, but for $l=6$.} \\label{Gf} \\end{figure}   With large redshift surveys becoming available in the next few years, the proposed statistical method appears as a promising tool for analyzing patterns in the galaxy distribution. In practice real galaxy catalogs have an incomplete sky coverage which implies, for the observed harmonic coefficients, non-zero mean values and a coupling to different modes of the all-sky coefficients. This occurs because of the  tensor window function of the angular mask arising from partial sky coverage. The required formalism  has already been applied in the literature to angular catalogs (Scharf et al. 1992; Scharf et al. 1993) and its application to the $Q_l$ coefficient is straightforward. Further complications arising from the radial selection function of the catalog and redshift space distortions have also been considered (Scharf et al. 1993; Ballinger, Heavens \\& Taylor 1995). For the functions $G_l(r)$ the analysis is much more cumbersome and is beyond the scope of this paper, which is intended to introduce the method and to illustrate its features by applications to numerical simulations.  An issue which has not been considered is whether this statistical tool can successfully be used to check the consistency of clustering data with cosmic models of structure formation. The analysis performed reveals (Fig. \\ref{Gd}) that any method which quantifies in a sensible way the presence of filaments or  planar structures in the cosmic web is also prone to cosmic variance. When different scales or levels of bias are considered, then the fluctuations are reduced (Figs. \\ref{Ge} \\& \\ref{Gf}). The results obtained suggest that an application of the estimator $G_l(r)$ to real data will require different scales to be probed, in order to give useful constraints on cosmic theories.   \\noindent {\\bf Acknowledgements} The author is grateful to the referee (M.Kerscher) for useful comments which improved the article."
    },
    "0011/astro-ph0011121_arXiv.txt": {
        "abstract": "We show how the simultaneous presence of a strong quasi periodic oscillation (QPO) of period $\\sim$10 seconds in the optical and X-ray lightcurves of the X-ray transient XTE J1118+480 can be used to obtain information about the nature of the accretion flow around the source. The unusually high optical-to-X-ray flux ratio and the QPO observed simultaneously in both energy bands suggest that a significant fraction of the optical flux might originate close to the central source, where most of the X-rays are produced and be indicative of a magnetically dominated corona. We also show how the temporal evolution of the QPO can provide us with information on both the inner radius and the viscous properties of the optically thick accretion disc. ",
        "introduction": "The transient X-ray source XTE J1118+480 was discovered with the RXTE All-Sky Monitor on March 29, 2000 (Remillard et al. 2000). At the peak of its outburst, it showed an X-ray spectrum typical of black hole candidates (BHC) in their {\\it hard} state, with a prominent power-law component (photon index of about 1.8). The 13th magnitude optical counterpart, discovered by Uemura, Kato \\& Yamaoka (2000), exhibited a spectrum fairly typical of X-ray Novae in outburst (Dubus et al. 2000 and references therein). A strong QPO with frequency $\\nu_{\\rm QPO} \\sim 0.1$ Hz has been found in the X-ray power density spectrum (PDS) of the source (Wood et al. 2000); more surprisingly, the optical lightcurves also showed a prominent QPO at the same frequency (Haswell et al. 2000). The QPO frequency drifted to higher values systematically in both bands for about two months.  As discussed in detail in Merloni, Di Metteo \\& Fabian (2000) (MDF), the high optical-to-X-ray flux ratio and the optical/UV variability argues for the presence of significant cyclo-synchrotron (CS) radiation from active regions in a magnetically dominated corona in this source. Assuming that a fraction of the observed optical flux corresponding to its oscillating  part ($\\sim 30\\%$) is due to CS emission from a magnetized corona, we modeled the spectrum taking into account reprocessing of coronal radiation in the accretion disc (according to the height of the active regions) and all the relevant radiative processes (see MDF for details). The disc was assumed to extend down to the innermost stable orbit of a non-rotating black hole ($R_{\\rm in}=6GM/c^2=3R_{\\rm S}$). We found that, in order to match the high (compared to X-rays) optical flux, both a low accretion rate ($\\dot M/\\dot M_{\\rm Edd}=\\dot m\\simeq 0.01$), and a high fraction of accretion power dissipated in the corona ($f>0.9$) are needed; also the active regions need to be high above the disc (at about $10 R_{\\rm S}$)\\footnote{ Alternatively, a good agreement with the data could be achieved if the active regions move relativistically away from the disc (MDF, see also Markoff, Falcke \\& Fender 2000 for a jet model of the SED)}.  Hynes et al. (2000) reporting on {\\it EUVE} observations of the source showed that the low measured EUV flux is inconsistent with an optically thick disc extending down to radii $R < 1000 R_{\\rm S}$ (but see Dubus et al. 2000 for a different estimate of the EUV flux). Based on this, they favor an inner advection-dominated flow (ADAF) surrounded by a geometrically thin disc truncated at hundreds of $R_{\\rm S}$. In the next section we show how to use the variability information to put constraints on the inner disc location. ",
        "conclusions": "The main rationale for our model is the simultaneous presence of a strong QPO in the optical and X-ray lightcurves of XTE J1118+480, suggesting that the fluxes in the two bands both originate from the same region in the inner part of the accretion flow. Self-absorbed CS emission is the natural candidate to explain the optical variability. Such emission is expected in any magnetic corona model, and the inferred magnetic field value ($B\\approx 2 \\times 10^6$ G) is the one predicted to arise when the source is in the {\\it hard} state (MDF). We have also shown how the evolution of the QPO can be satisfactory accounted for if the QPO is produced  by a modulation whose frequency is equal to twice the relativistic nodal precession in a radiation pressure dominated disc and the location of such modulation (that is perhaps where the inner accretion disc undergoes a transition and is {\\it at most} a few tens of Schwarzschild radii) viscously drift inwards."
    },
    "0011/astro-ph0011067_arXiv.txt": {
        "abstract": "\\baselineskip15pt \\mbox{}\\\\ Cosmological gamma-ray bursts are probably powered by systems harboring a rotating black hole. We show that frame-dragging creates baryon-poor outflows in a differentially rotating gap along an open magnetic flux-tube with dissipation ${P}=\\Omega_T(\\Omega_H-2\\Omega_T)A_\\phi^2$, where $2\\pi A_\\phi$ denotes the flux in the open tube, for angular velocities $\\Omega_H$ and $\\Omega_T$ of the black hole and the torus, respectively. The output in photon-pair ouflow will be reprocessed in an intervening hypernova progenitor wind or the ISM.  \\mbox{}\\\\ {\\em Subject headings}: gamma-rays: bursts, theory \\mbox{}\\\\ ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011073_arXiv.txt": {
        "abstract": "The Red-Sequence Cluster Survey (RCS) is a new galaxy cluster survey designed to provide a large sample of optically selected 0.1$<$z$<$1.4 clusters. The planned survey data are 100 square degrees of two color ($R$ and $z'$) imaging, with a 5$\\sigma$ depth $\\sim$2 mag past M$^{*}$ at $z=1$. The combined depth and area of the RCS make it the widest field, moderately deep survey ever undertaken using 4m class telescopes.  This paper gives a brief outline of the RCS survey, with particular emphasis on the data reduction strategy.  The remainder of the paper focuses on preliminary results from the first set of completely reduced data ($\\sim$10 deg$^2$, of the $\\sim$60 deg$^2$ in hand). We provide a new example of a rich $z>1$ cluster, illustrative of the dozens discovered in the data so far. Some of the possible science to come from the RCS is illustrated by a qualitative indication of $\\Omega_{M}$ from the first 1/10th of the survey data. A high-precision measurement of the 2-point correlation function of luminous early-type galaxies at $0.4<z<1.2$ is also shown. ",
        "introduction": "\\paragraph{} Galaxy clusters are the largest gravitationally bound structures in the universe, and as such are markedly important in cosmology and studies of galaxy evolution. Much of the work on clusters done to date at low to intermediate redshifts (typically limited to $\\sim0.8$ or so) would benefit greatly from extension to higher redshifts. However, the cluster samples required for such studies do not exist. The Red-Sequence Cluster Survey (RCS) is designed to rectify this problem. The survey is designed to find galaxy clusters in large numbers out to $z\\sim1.4$, spread over a wide range of RA and DEC in order to facilitate follow-up.  \\paragraph{} The first observations for the RCS were taken in May 1999, at the CFHT, and at the time of writing we have acquired $\\sim60\\%$ of the 100 square degrees of planned survey data. Data acquisition will be completed by the end of 2001. The survey data consist of 22 survey patches (ten $2.1\\times2.3$ deg$^2$ patches from the CFHT, and twelve $1.8\\times2.4$ deg$^2$ patches from the CTIO 4m telescope) imaged in the $R_C$ and $z'$ filters to $5\\sigma$ point source depths of 25.2 and 23.6 mags respectively. This depth has been chosen to facilitate finding $z\\sim1$ galaxy clusters, which are located in the survey data using an algorithm based on detecting the early-type galaxies in the clusters cores (Gladders \\& Yee 2000a). The RCS has been described in more detail elsewhere (Gladders \\& Yee 2000b), and we will not provide a repetition here. The interested reader is referred to the survey homepage, {\\texttt {http://www.astro.utoronto.ca/$\\sim$gladders/RCS/}}, which contains a description of the survey strategy, goals, and other relevant data and publications. \\paragraph{} The purpose of this paper, rather, is to provide some of the details of the survey design strategy, as these are germane to the topic of this meeting. In addition, we provide some examples of new results coming from the RCS, including a preliminary indication of $\\Omega_M$, an example of $z\\sim1$ galaxy clusters, and some of the secondary science related to the identification of individual early-type galaxies using photometric techniques over an unprecedented volume to $z\\sim1$. ",
        "conclusions": ""
    },
    "0011/astro-ph0011245_arXiv.txt": {
        "abstract": "{ Recent {\\em XMM-Newton} observations reveal an extended ($r \\approx 150''$) low-surface brightness X-ray halo in the supernova remnant G21.5-0.9. The near circular symmetry, the lack of any limb brightening and the non-thermal spectral form, all favour an interpretation of this outer halo as an extension of the central synchrotron nebula rather than as a shell formed by the supernova blast wave and ejecta. The X-ray spectrum of the nebula exhibits a marked spectral softening with radius, with the power-law spectral index varying from $\\Gamma = 1.63 \\pm 0.04$ in the core to $\\Gamma = 2.45 \\pm 0.06$  at the edge of the halo. Similar spectral trends are seen in other Crab-like remnants and reflect the impact of the synchrotron radiation losses on very high energy electrons as they diffuse out from the inner nebula. A preliminary timing analysis provides no evidence for any pulsed X-ray emission from the core of G21.5-0.9. ",
        "introduction": "\\label{sec:intr}  Of the $\\sim 225$ known supernova remnants (SNRs) in our galaxy no more than $5\\%$ are classified as Crab-like remnants (Green \\cite{green00}). Prominent members of this class of SNR, also known as plerions, include the Crab Nebula, CTB 87, and 3C58. The characteristic centre-filled radio and X-ray morphologies of the Crab-like SNRs is thought to be due to the presence of an active pulsar which powers a bright synchrotron nebula. In the case of the Crab nebula, 33 ms pulsations testify to the presence of the central spinning neutron star, although in many Crab-like systems there is no direct evidence for pulsed emission.  Typically a Crab-like SNR exhibits a flat power-law spectrum in the radio regime which eventually steepens at shorter wavelengths to join smoothly to a similarly featureless power-law X-ray continuum. This non-thermal form for the radio to X-ray spectrum is a consequence of the continuous injection by the pulsar of high energy electrons, which suffer radiation and adiabatic losses as they diffuse through the nebula; however the details of the process are undoubtedly complex (e.g. Reynolds \\& Chanan \\cite{reynolds84}). A feature of the Crab-like remnants, which distinguishes them from both shell-like and composite-type SNRs (the latter typically exhibiting a filled-centre X-ray morphology within a radio shell), is the lack of any evidence for an outer shell structure marking the progress of the blastwave from the original supernova explosion.  G21.5-0.9 is a SNR with many of the characteristics of a Crab-like remnant (e.g. Becker \\& Szymkowiak \\cite{becker81}; Asaoka \\& Koyama \\cite{asaoka90}). To date there has been no detection of pulsed emission in either the radio or X-ray bands (e.g. Biggs \\& Lyne \\cite{biggs96}; Slane et al. \\cite{slane00}). However, recent X-ray observations made by {\\em Chandra} have pinpointed the probable location of the pulsar, namely a very compact central core on a scale of $\\sim 2''$ embedded within a more extended ($\\sim 30''$ radius) synchrotron nebula (Slane et al. \\cite{slane00}). A faint extended halo was also detected which might correspond to the outer ``shell'' formed from the expanding ejecta and the passage of the supernova-driven blastwave (Slane et al. \\cite{slane00}). In this interpretation G21.5-0.9 would be best described as a composite remnant. Here we report further recent X-ray observations of G21.5-0.9 made by {\\em XMM-Newton} which help to clarify the nature of the extended X-ray halo in this object.   \\begin{figure*} \\centering \\parbox{17.0cm}{\\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{XMM38_f1.eps}}}} \\hfill \\caption[]{{\\it Left panel:} The raw {\\em XMM-Newton} EPIC PN image of G21.5-0.9. The pixel size is $4''$ and the image is $8'$ on a side. A logarithmic intensity scaling has been applied. The vertical and horizontal white bands correspond to the gaps between adjacent CCD chips of the PN camera. {\\it Right panel:} EPIC MOS image of G21.5-0.9. The pixel size is $1''$ and the image is $8'$ on a side.  The image has been spatial filtered with a Gaussian smoothing mask with width $\\sigma = 2 $ pixel. Again logarithmic intensity scaling has been applied to emphasise low surface-brightness features. } \\label{colour} \\end{figure*} ",
        "conclusions": "In a recent paper reporting Chandra observations, Slane et al. (\\cite{slane00}) consider the possibility that the low-surface brightness halo surrounding the central synchrotron nebula in G21.5-0.9 might be the shell formed from the ejecta and blastwave of the original supernova explosion.  The present {\\em XMM-Newton} observations help delineate the spatial extent and morphology of this component. Crucially the new spectral information demonstrates that the extended halo has a spectral form devoid of any significant line features. If the halo emission is thermal then the lack of line emission implies that the plasma is far from ionization equilibrium. Our NEI modelling confirms a low ionization state with $n_{\\rm e}t \\approx 3 \\times 10^{8} \\rm~cm^{-3}~s$. The temperature derived for the continuum (bremsstrahlung) emission is $4 - 5$ keV which is rather hot, even for a very young SNR. Assuming the distance to G21.5-0.9 is 5 kpc (Davelaar, Smith \\& Becker \\cite{davelaar86}) and taking a maximum shell expansion velocity of $10,000 \\rm~km~s^{-1}$, we can estimate the (minimum) elapse time since the shock heating of the bulk of the X-ray emitting gas in the outer halo to be $\\sim 100$ yrs. Applying the above constraint on $n_{\\rm e}t$ then gives $n_{\\rm e} < 0.1 \\rm~cm^{-3}$. However, the observed X-ray luminosity of the outer halo, if interpreted as bremsstrallung radiation, sets a requirement for a significantly higher electron density, $n_{\\rm e} \\approx 0.5 \\rm~cm^{-3}$. There would appear therefore to be considerable difficulties in interpreting the extended X-ray halo in G21.5-0.9 as a thermal shell.  There are known examples of SNR shells which exhibit non-thermal X-ray spectra due to local particle acceleration (e.g. SN 1006, Koyama et al. \\cite{koyama95}; G347.3-0.5, Slane et al. \\cite{slane99}). However, in the case of G21.5-0.9 the lack of any limb brightening, the smooth transition in spectral index throughout the remnant and the remarkable circular symmetry, all suggest an interpretation of the outer halo as an extension of the central synchrotron nebula. We conclude therefore, on the basis of the present observations, that G21.5-0.9 is a true Crab-like system rather than a composite (centre-filled plus thermal shell) object.  The observed spectral softening with radius in G21.5-0.9 most likely reflects the impact of synchrotron radiation losses on very high energy electrons as they flow from the region of the termination shock around the pulsar to the edges of the nebula. For example, if the nebula magnetic field is  $\\sim 0.4 $ mG (Slane et al. \\cite{slane00} ), then the synchrotron lifetime of the $\\sim 20$ TeV electrons producing 5 keV X-rays is only $\\sim 2$ yrs. Similar spectral trends due to ``synchrotron burn-off'' are seen in other Crab-like SNR, such as 3C58 (Torii et al. \\cite{torii00}). In fact the variation in photon spectral index from a value of 1.63 at the centre of the G21.5-0.9 to 2.45 at the edge of its halo is amazingly similar to the index range seen in the Crab nebula (Willingale et al. \\cite{willingale01}). In the case of the Crab nebula the spatial and spectral distribution of the optical to hard X-ray continuum has been explained in terms a magnetohydrodynamic flow of plasma and magnetic flux in a particle dominated wind (e.g. Kennel \\& Coroniti \\cite{kennel84}). For G21.5-0.9, the origin of halo which represents a ``plateau'' in the surface brightness extending from $r \\approx 50'' - 130''$ remains unclear, but could possibly be the result of an unusual magnetic field geometry (see below).  One interesting contrast is that the radio extent of the Crab nebula is roughly a factor four greater than that observed in X-rays, whereas for G21.5-0.9 the dimensions of the X-ray halo exceeds that of the radio synchrotron nebula (which has a major axis of $\\sim 90''$; F{\\\"u}rst et al. \\cite{furst88}) by a similar factor. Slane et al. (\\cite{slane00}) note that no radio ``shell'' is detectable to a (1$\\sigma$) threshold of $4 \\times 10^{-21} \\rm~W~m^{-2}~Hz^{-2}~sr^{-1}$ at 1 GHz. If we make the assumption that the ratio of the core to halo flux is the same in the X-ray and radio regimes, then we predict the average surface brightness of an underlying radio halo (between $r = 60'' - 130 ''$) to be roughly 5 times this threshold.  The inference is that very deep radio observations may well reveal the presence of an extended radio halo in G21.5-0.9.  The differences in the spatial morphology of the Crab and G21.5-0.9 could be due to the viewing orientation. Specifically the near circular symmetry of G21.5-0.9 might be explained if our line of sight is reasonably aligned with the spin axis of the resident pulsar, so as to give a near face-on view of a putative inner torus.  This hypothesis might also explain the non-detection of pulsed X-ray emission from the core of G21.5-0.9, but is hard to reconcile with the axisymmetric structure seen in high frequency radio maps and the near radial distribution of the magnetic field lines (F{\\\"u}rst et al. \\cite{furst88}). In this setting the spur-like feature which emanates northwards from the core could a ``plerionic'' wisp associated with a complex inner magnetic field and/or jet structure."
    },
    "0011/astro-ph0011135_arXiv.txt": {
        "abstract": "This paper presents diffraction-limited 1$-$18$\\mu$m images of the young quadruple star system HD 98800 obtained with the W. M. Keck 10-m telescopes using speckle and adaptive optics imaging at near-infrared wavelengths and direct imaging at mid-infrared wavelengths. The two components of the visual binary, A and B, both themselves spectroscopic binaries, were separable at all wavelengths, allowing us to determine their stellar and circumstellar properties.  Combining these observations with spectroscopic data from the literature, we derive an age of $\\sim$10$^7$ years, masses of 0.93 and 0.64 M$_{\\sun}$ and an inclination angle of 58$^{\\circ}$ for the spectroscopic components of HD 98800 B, and an age of $\\sim$10$^7$ years and a mass of 1.1 M$_{\\sun}$ for HD 98800 Aa.  Our data confirm that the large mid-infrared excess is entirely associated with HD 98800 B.  This excess exhibits a black body temperature of 150 K and a strong 10$\\mu$m silicate emission feature.  The theoretical equilibrium radius of large, perfectly absorbing, 150 K grains around HD 98800 B is 2.4 AU, suggesting a circum-spectroscopic binary distribution. Our observations set important upper limits on the size of the inner dust radius of $\\sim$2 AU (from the mid-infrared data) and on the quantity of scattered light of $<$10\\% (from the H-band data).  For an inner radius of 2 AU, the dust distribution must have a height of at least 1 AU to account for the fractional dust luminosity of $\\sim$20\\% L$_B$. Based on the scattered light limit, the dust grains responsible for the excess emission must have an albedo of $<$0.33. The presence of the prominent silicate emission feature at 10$\\mu$m implies dust grain radii of $\\ga$2$\\mu$m. The total mass of the dust, located in a circumbinary disk around the HD 98800 B, is $>$0.002M$_{\\earth}$. The orbital dynamics of the A$-$B pair are probably responsible for the unusual disk geometry. ",
        "introduction": "HD 98800 is one of the most frequently studied members of the TW Hydra Association because of its notably large fractional infrared (IR) luminosity excess, first detected by the IRAS satellite \\citep{wal88, zuc93}. Hipparcos results reported by \\citet{fav98} indicate that HD 98800 is at a distance of 47.6 pc, which has led to recent age estimates for this K5 system of $\\sim$7 Myr, suggesting that the stars are pre$-$main-sequence (PMS) (Fekel and Bopp 1993; Soderblom et al. 1996; Webb et al. 1999; Jensen, Cohen, and Neuh\\\"auser 1998). The lack of obvious signatures of active accretion, such as hydrogen emission lines \\citep{sod96, web99} and near-IR excess (e.g. Low, Hines, and Schneider 1999), implies that the observed mid-IR excess probably originates in a dusty disk with a central gap.  Complicating both the dynamics and the analysis of the HD 98800 system is its multiplicity.  \\citet{inn09} identified it as a visual binary; its current north-south angular separation of 0.$''$8 corresponds to a projected linear separation of 38 AU.  Orbital solutions for the visual binary are not unique because only linear motion has been observed during the last nine decades. \\citet{tor95} initially found a family of orbits with extremely high eccentricity ($\\sim$0.9), e, and periods, P, $>$10$^5$ years. Using the Hipparcos data, and with knowledge of the PMS nature of the system, \\citet{tok99} published a different set of orbital solutions with moderate 300$-$430 year periods and 0.3$-$0.6 eccentricities.  Tokovinin's orbits are probably more reliable since they incorporate better values for the component masses and the Hipparcos distance, thus improving the determination of e and P \\citep{tor00}. Spectroscopic measurements have shown that each component of the visual binary has a companion. The southern component, HD 98800 A, is a single-lined spectroscopic binary (SB) with P$=$262 days and e$=$0.484, while the northern component, HD 98800 B, is a double-lined SB with P$=$315 days and e$=$0.781 \\citep{tor95}. The presence of numerous components raises questions regarding the relationship between the stars and the dust: Where does the dust reside?  Is the interaction between the multiple stellar components and the dust responsible for HD 98800's unusually large fractional IR excess?  And more speculatively, can planets form and survive within this system?  To address these questions, we carried out a 1$-$18$\\mu$m diffraction-limited imaging study with the W. M. Keck 10-m telescopes, complementing results presented in several recent papers. \\citet{sod98} and \\citet{low99} used, respectively, the WFPC2 and NICMOS cameras on the {\\it Hubble Space Telescope} to characterize the stellar parameters of the individual A and B components and to search for scattered light from the dust.  \\citet{geh99} used the ESO 3.6-m telescope to study the system at 4.7$\\mu$m and 9.7$\\mu$m, locating the bulk of the dust around the B component.  \\citet{koe00} combined mid-IR with published millimeter wavelength data, dramatically illustrating that the excess emission is associated with HD 98800 B and modelling the dust distribution. With a factor of 2$-$4 times higher spatial resolution than most previous observations, the measurements in this study are used to characterize the physical properties and spatial distribution of the dust and stars in HD 98800 and to construct a paradigm for the system. The observations are described in \\S2 and the analysis and results in \\S3. A discussion is provided in \\S4 and a summary in \\S5. ",
        "conclusions": "\\subsection{Disk Structure}  The results presented in \\S 3 and in the literature constrain the structure of the disk around HD 98800 B. Our estimates for the albedo, $<$0.33, dust temperature, 150 K, and the total stellar luminosity, 0.58L$_{\\sun}$, allow us to calculate the theoretical equilibrium radius of the dust.  Because the luminosity comes from two sources, the stars in the highly eccentric spectroscopic binary, we make the simplifying assumption here that the two are equally bright.  A numerical solution for the time averaged stellar luminosity as seen from a disk gives a dust equilibrium radius of $\\sim$2.4 AU from the center of mass of the system.  The criterion of \\citet{art94}, [(1/2)$a$(1$+e$)]$+a$, one half the apastron separation of the stars plus the semi-major axis, $a$, can be used to determine the inner radius of a stable, tidally truncated circumbinary disk. Combining the data given by \\citet{tor95} for $a_1$sin$i$, $a_2$sin$i$ and the eccentricity, $e$, with our estimate of $i=$58$^{\\circ}$, yields $a=$1.07 AU; for a stable disk, the inner radius must therefore be $\\ga$2 AU. The data presented in Table 1 indicate that the 3$\\sigma$ upper observational limit on the inner radius is between 1.8 and 2.1 AU (\\S3.1), barely consistent with this result.  From the slight photometric variability of the HD 98800 system seen in the Hipparchos data \\citep{sod98}, \\citet{tok99} infers that the Ba$-$Bb binary is seen at an inclination of $\\sim$90$^{\\circ}$ through the circumbinary dust. Given the 10$^7$ year age of HD 98800 B, and the 315 day orbital period of Ba$-$Bb, the dust distribution and the stellar orbit are probably coplanar (Lubow 2000). Thus, the 58$^{\\circ}$ inclination angle derived for the spectroscopic binary also applies to the dust.  The large fractional luminosity, 0.2L$_B$, requires that the dust be vertically extended, since, at a radius of 2 AU, a thin disk intercepts far too little starlight to account for the dust brightness.  A disk of perfectly absorbing grains would have to have a height of at least 1 AU in order to account for the large fractional luminosity if all the light reradiated by the disk were intercepted by the grains at a radius of $\\sim$2 AU. Given the upper bound on the albedo, 0.33, the height of the disk must be between 1 and 1.5 AU. For a spherical distribution of dust, we expect the fractional luminosity to be greater than the 0.2L$_B$ observed; thus, a thick disk geometry is most likely.  The dust around the HD 98800 B spectroscopic binary is located in a circumbinary disk with an equilibrium radius of 2.4 AU, an inner gap of $\\sim$2 AU, and a vertical extent of $>$1 AU.  To account for the observed 8.8$\\mu$m excess flux density of 0.61 Jy, the outer radius of a disk at 47.6 pc with a 2 AU inner gap and T$=$150 K must be $\\sim$5 AU if the origin of the emission is primarily black body. (The value listed in Table 2 for the 8.8$\\mu$m flux density includes the stellar photosphere; Figure 6 shows that the dust to photosphere brightness ratio is 2.1 at 8.8$\\mu$m.) Cooler dust probably extends further out than 5 AU \\citep{koe00}, however, it will eventually be truncated by the dynamical action of the visual binary's orbit at about one third of the periastron distance between HD 98800 A and B, i.e. at least $\\sim$10$-$15 AU, depending on projection effects.  \\subsection{A Paradigm for the HD 98800 System}  We believe that the orbit of the visual binary, HD 98800 A$-$B, is responsible for the current configuration of the system, since the inclination of this wide pair is 87$^{\\circ}-$89$^{\\circ}$ \\citep{tor95, tok99}, while the inclination of HD 98800 Ba$-$Bb and the associated dust disk is $\\sim$58$^{\\circ}$. The A$-$B and the Ba$-$Bb orbits, and therefore the A$-$B orbit and the B-component dust disk, are not coplanar. Tidal forces not only may have ripped away circumbinary material from HD 98800 A in the past, but also may perturb the HD 98800 B circumbinary disk in the present, maintaining the small grain population against Poynting-Robertson drag and giving rise to its unusual vertical extension. Disk warping as a result of binary perturbations from a non-coplanar companion has been described by \\citet{ter96} and \\citet{lar96}. These perturbations could cause precession of the HD 98800 B dust orbital plane, maintaining the 58$^{\\circ}$ inclination constant but varying the orientation of the disk.  Such changes would cause the disk to be at times perpendicular and at times parallel to the A$-$B orbital plane. We speculate that similar dynamics destroyed the disk around HD 98800 A, anomalous in that no evidence of dust is apparent, and currently maintain the configuration of the HD 98800 B dust, still prominent although the stellar ages are about 400 times greater than the Poynting-Robertson lifetime of the grains, and apparently significantly vertically extended. Orbital periods in the models proposed by \\citet{tok99} range from 300 to 430 years, allowing for relatively short term periodic perturbations of the HD 98800 B disk, initiated during periastron passage with HD 98800 A."
    },
    "0011/astro-ph0011303_arXiv.txt": {
        "abstract": "Unusual patterns of recombination emission from gas ionized by metal-free stars may distinguish early star-forming galaxies from their present-day counterparts.  This pattern arises from the harder ionizing spectrum expected from metal-free stars, which strongly enhances the strength of \\ion{He}{2} recombination lines. Our calculations indicate that line fluxes of \\ion{He}{2} $\\lambda$1640 and $\\lambda$4686 are sufficiently large to be detected by narrowband and spectroscopic searches for high-redshift emission-line sources at $z \\sim 5$ using current instruments. An unknown fraction of Ly$\\alpha$ emitters may harbor low-metallicity or metal-free stars.  As the predicted \\ion{He}{2} $\\lambda$1640 flux is comparable to and may exceed hydrogen Ly$\\alpha$, searches for high-redshift galaxies should consider \\ion{He}{2} recombination lines as possible identifications for single emission lines in observed spectra.  Spectra of metal-free stars may show both \\ion{H}{1} and \\ion{He}{2} emission lines, improving the constraints on their redshift and identification.  We assess the considerable uncertainties that affect our expectations for the detection and identification of true first-generation stars with present search techniques, including the role of stellar mass loss in spectral evolution and the confusion of ionization by primordial stellar sources and AGN in the early universe. ",
        "introduction": "The first stars have long attracted interest from a cosmological point of view, and recent detailed studies of their properties have renewed interest in their cosmological importance.  Within the constraints of big-bang nucleosynthesis, these stars must have been metal free, which modifies their structure in fundamental ways:  metal free stars have smaller radii, hotter cores, and higher effective temperatures than their counterparts of equal mass but finite metallicity (Tumlinson \\& Shull 2000). These stars are leading candidates for the sources that reionized the intergalactic medium (Gnedin 2000; Haiman \\& Loeb 1997), and they produced the first heavy elements that enriched subsequent stellar generations.  Tumlinson \\& Shull (2000) found that these stars, here called Population III or simply metal-free, have harder spectra and emit more ionizing photons in the \\ion{He}{2} ionizing continuum than do stars of typical solar or sub-solar metallicity.  Thus, they may dominate \\ion{He}{2} reionization in the high-redshift universe.  Distinctive spectral features of their surrounding nebulae may provide a means of detecting the first stars, during and after the epoch of reionization.  Tumlinson \\& Shull (2000) noted that the hard spectra of Pop III stars produce large \\ion{He}{3} regions, which may emit detectable \\ion{He}{2} recombination emission.  Recombination lines of \\ion{He}{2} at $\\lambda$1640 ($n = 3 \\rightarrow$ 2), $\\lambda$3203 ($n = 5 \\rightarrow$ 3), and $\\lambda$4686 ($n = 4 \\rightarrow$ 3) are particularly attractive for this purpose, because they suffer minimal effects of scattering by gas and decreasing attenuation by intervening dust.  Thus, an assessment of the near-term prospects for discovering metal-free stellar sources is timely.  However, uncertainties about the primeval initial mass function (IMF), the importance of stellar mass loss at zero metallicity, and the features of stellar evolution must be addressed.  In this {\\em Letter}, we consider the possibility of detecting metal-free stars using \\ion{He}{2} recombination emission, and the uncertainties inherent in this technique.  Main-sequence mass loss, not well understood from an evolutionary standpoint, can affect the detectability of the first stars via nebular emission lines.  Emission line ratios and line profiles may discriminate between stars and active galactic nuclei (AGN) as sources of ionization.  We address these issues in the context of existing tracks and assess their importance to the detectability of metal-free stars.  In \\S~2 we derive an expression for the observed \\ion{He}{2} emission-line flux of a metal-free stellar population in terms of the star formation rate and evolutionary properties. In \\S~3 we use evolutionary tracks and model atmospheres to assess the effects of stellar evolution on the detectability of metal-free stellar populations.  In \\S~4 we compare our predictions to current observational searches for high-redshift stellar populations, including present-day emission-line techniques, and we discuss the ambiguity between emission lines produced by metal-free stars and AGN.  In \\S~5 we summarize our results and comment on directions for future work. ",
        "conclusions": "In summary, our main conclusions are the following:  \\begin{itemize} \\item[1.] Predicted fluxes for \\ion{He}{2} emission lines from nebulae ionized by metal-free stars are comparable to the detection limits of present-day narrowband and spectroscopic searches for high-redshift galaxies. Observational searches for these stars at $z \\sim 3 - 5$ are feasible now at optical wavelengths.  Higher redshift surveys may need to await space-based infrared instruments.  \\item[2.] Mass loss by metal-free stars may restrict or enhance their detectability with emission-line techniques.  Careful calibration with detailed evolutionary tracks and theoretical work on mass loss from metal-free stars are necessary to refine this program.  \\item[3.] If multiple line detections are made for a single object, \\ion{He}{2} and \\ion{H}{1} emission lines may provide important diagnostics of the stellar parameters. These line ratios may provide a means of distinguishing stellar and AGN sources even if the line profiles are unresolved.  \\end{itemize}"
    },
    "0011/astro-ph0011186_arXiv.txt": {
        "abstract": "The Virgo and A1795 clusters of galaxies were re-observed by EUVE with {\\it in situ} background measurements by pointing at small offsets.  Earlier, a similar re-observational strategy applied to the cluster A2199 revealed that the background radial profile was consistent with a flat distribution,  and therefore the original method of extracting cluster EUV signals by the subtraction of an aymptotically determined background was valid. It is shown here that the same conclusions hold for the current sample. A model of the background was obtained from its known properties and the {\\it in situ} measurements, and the subtracted cluster fluxes remain in agreement with those reported in our discovery papers.  They are also consistent with results from the most conservative procedure of direct point-to-point subtraction of the {\\it in situ} background and proper error propagation, which still preserves the existence of the EUV excess and its rising radial trend. We present evidence which argues against the soft excess as due to peculiarities in the line-of-sight Galactic absorption. The data appear to favor a thermal origin of the emission. ",
        "introduction": "Clusters of galaxies are sources which emit EUV and soft X--ray radiation ($\\sim$ 0.1-- 0.4 keV) substantially in excess of that expected from the hot intra--cluster medium (ICM) at X--ray temperatures (a few keV). The {\\it Extreme Ultraviolet Explorer} mission (Bowyer and Malina 1991) was pivotal to the discovery of this phenomenon, as it opened a unique window to the EUV sky with the Deep Survey (DS) Lex/B filter ($\\sim$ 65--190 eV, hereafter the EUV band). Cluster soft--excess (CSE) emission was reported for the Virgo (Lieu et al. 1996a), Coma (Lieu et al. 1996b), A1795 (Mittaz, Lieu and Lockman 1998) and A2199 clusters (Lieu, Bonamente and Mittaz 1999). Recent attributions of these findings to spatial variations of the background (Bowyer, Berghoefer and Korpela 1999) were based upon an incorrect way of analyzing the EUVE  data, and the decisive way of assessing the existence of CSE emission is that of an {\\it in situ} background measurement by pointing time contiguously at small offsets from the cluster. This approach,  applied to a recent EUVE  re-observation of A2199 (Lieu et al. 1999b), clearly demonstrated that the offset pointing provides an accurate background template for the cluster observation, whereas an arbitrary blank field does not.  With the aid of new EUVE observations, performed along with acquisition of {\\it in situ} background, we were able to confirm the previous detection of CSE even in the most conservative case when we assumed no {\\it a priori} understanding of the background distribution, so that a point-to-point subtraction of the background had to be adopted. We provide in this analysis two ways of utilizing the background data, one is to develop a model of the background and the other is direct subtraction of the offset field as in the case of A2199. ",
        "conclusions": ""
    },
    "0011/astro-ph0011465_arXiv.txt": {
        "abstract": "The morphology and kinematics of molecular clouds (MCs) are best explained as the consequence of super--sonic turbulence. Super--sonic turbulence fragments MCs into dense sheets, filaments and cores and large low density ``voids'', via the action of highly radiative shocks. We refer to this process as {\\it turbulent fragmentation}.  In this work we derive the mass distribution of gravitationally unstable cores generated by the process of turbulent fragmentation.  The mass distribution above one solar mass depends primarily on the power spectrum of the turbulent flow and on the jump conditions for isothermal shocks in a magnetized gas. For a power spectrum index $\\beta=1.74$, consistent with Larson's velocity dispersion--size relation as well as with new numerical and analytic results on super--sonic turbulence, we obtain a power law mass distribution of dense cores with a slope equal to $3/(4-\\beta) = 1.33$, consistent with the slope of the stellar IMF. Below one  solar mass, the mass distribution flattens and turns  around at a fraction of  a solar mass, as  observed for the stellar IMF in a number  of stellar clusters, because only the densest cores  are gravitationally  unstable.   The mass  distribution at  low masses  is  determined by  the  probability  distribution  of the  gas density, which is  known to  be approximately Log--Normal for an isothermal turbulent  gas.   The intermittent  nature  of  the turbulent  density distribution is thus responsible for  the existence of a significant number of small collapsing cores, even of sub--stellar mass.  Since  turbulent  fragmentation  is  unavoidable  in  super--sonically turbulent molecular clouds, and given the success of the present model in predicting the observed shape  of the stellar IMF, we conclude that turbulent fragmentation is essential to the origin of the stellar IMF. ",
        "introduction": "The process of star formation, particularly the origin of the stellar initial mass function (IMF), is a fundamental problem in astrophysics. Photometric properties and chemical evolution of galaxies depend on their stellar content. The process of galaxy formation cannot be described independently of the process of star formation, since galaxies are partly made of stars. Stars are also an important energy source for the interstellar medium of galaxies.  Stars are formed in molecular clouds (MCs), which have been the focus of the research on star formation for more than two decades. Currently, there is no generally accepted theory of star formation, capable of predicting the star formation rate and the stellar IMF based on the physical properties of MCs. This is hardly surprising, since turbulent motions are ubiquitously observed in MCs, and the physics of turbulence is poorly understood, due to the great mathematical complexity of the fluid equations. Magnetic field, self--gravity and high Mach numbers further increase the complexity.  The steady growth of computer performance has now made large three--dimensional numerical simulations of super--sonic magneto-hydrodynamic (MHD) turbulence feasible (Padoan \\& Nordlund 1997, 1999; Stone, Ostriker \\& Gammie 1998; Mac Low et al.\\ 1998; Padoan, Zweibel \\& Nordlund 2000; Klessen, Heitsch \\& Mac Low 2000; Mac Low \\& Ossenkopf 2000; Ostriker, Stone \\& Gammie 2000; Heitsch, Mac Low \\& Klessen 2000; Padoan et al.\\ 2001a,b). Comparisons of numerical experiments with observational data have shown that super--sonic turbulence can explain the morphology and kinematics of MCs, and the formation of dense cores, provided that the motions are also super-Alfv{\\'e}nic (Padoan, Jones \\& Nordlund 1997; Padoan et al.\\ 1998; Padoan \\& Nordlund 1997, 1999; Padoan et al.\\ 1999; Padoan, Rosolowsky \\& Goodman 2001; Padoan et al.\\ 2001a, b). We refer to this process of formation of dense cores in MCs by super--sonic turbulence as {\\it turbulent fragmentation}. Since protostars evolve from the collapse of gravitationally unstable cores in MCs, even the stellar IMF could then be the result of turbulent fragmentation, with the power law shape of the IMF ultimately being the consequence of the self--similar nature of turbulence. \\nocite{MacLow_Puebla98} \\nocite{Stone+98} \\nocite{Kimura+Tosa93} \\nocite{Lee+91} \\nocite{Vazquez-Semadeni+96} \\nocite{Klessen+2000} \\nocite{Padoan+97ext} \\nocite{Padoan+98per} \\nocite{Padoan+Nordlund98MHD} \\nocite{Padoan+98cat} \\nocite{Padoan+Nordlund97MHD} \\nocite{Ostriker+99} \\nocite{Passot+95} \\nocite{MacLow+98} \\nocite{MacLow99} \\nocite{Padoan+2001frag} \\nocite{MacLow+Ossenkopf2000} \\nocite{Ostriker+2000} \\nocite{Heitsch+2000} \\nocite{Padoan+2000ad}  In this paper we do not use these increasingly sophisticated numerical simulations directly, but instead develop an analytic model.  The assumptions of the model are {\\em inspired} by the qualitative properties of the numerical model, but the current work does not {\\em depend} on any particular set of numerical models or results.  Previous analytic models by Larson (1992), Henriksen (1986, 1991) and Elmegreen (1993, 1997, 1999, 2000b) have derived the stellar IMF on the basis of the self--similar structure of MCs. Larson, assuming one dimensional accretion, predicted a rather steep IMF slope, equal to the MC fractal dimension, while Henriksen found the IMF slope to depend on both the MC fractal dimension and the relation between density and linear size for structures inside MCs. Elmegreen pointed out that the IMF that results from random sampling of a self--similar cloud has an exponent $x=1$ (Salpeter $x=1.35$), independent of the cloud fractal dimension. In Elmegreen's model the IMF is steeper than $x=1$ because the random sampling rate is assumed to be proportional to the square root of density, and because of ``mass competition''.  \\nocite{Larson92}  \\nocite{Henriksen86}  \\nocite{Henriksen91} \\nocite{Elmegreen97imf} \\nocite{Elmegreen99imf}  \\nocite{Elmegreen2000imf}  These models are based on assumptions about the cloud geometry justified by the apparent fractal structure of MCs (Beech 1987; Bazell \\& D\\'{e}sert 1988, Scalo 1990; Dickman, Horvath \\& Margulis 1990; Falgarone, Phillips \\& Walker 1991; Zimmermann, Stutzki \\& Winnewisser 1992; Henriksen 1991; Hetem \\& Lepine 1993; Vogelaar \\& Wakker 1994; Elmegreen \\& Falgarone 1996), but the processes responsible for generating the assumed geometry are not discussed in detail.  \\nocite{Beech87}  \\nocite{Bazell+Desert88}  \\nocite{Scalo90} \\nocite{Dickman+90}  \\nocite{Falgarone+91} \\nocite{Zimmermann+92} \\nocite{Henriksen91} \\nocite{Hetem+Lepine93} \\nocite{Vogelaar+Wakker94} \\nocite{Elmegreen+Falgarone96}  In a previous attempt to relate the stellar IMF to the physical properties of super--sonic turbulence (Padoan, Nordlund \\& Jones 1997), we obtained the distribution of the local Jeans' mass in super--sonic isothermal turbulence, from the probability distribution of the gas density. That work also provides a prediction for the stellar IMF, by identifying each local Jeans' mass with a protostar. Although the prediction of the lower mass cutoff and the low mass portion of the IMF might be roughly correct, this model under--estimates the number of massive stars relative to low mass stars (as any other Log-Normal IMF). The main reason, as pointed out by Scalo et al. (1998), is the unphysical assumption that the most massive stars originate from gas at relatively low density including only a small fraction of the total mass.  \\nocite{Padoan+97imf}  \\nocite{Scalo+98}  In the present work we consider specific properties of MC turbulence and their consequences for the formation of protostellar cores. In particular, we  i) assume approximate self--similarity of the super--sonic and super--Alfv{\\'e}nic velocity field, expressed by a power law shape of the power spectrum of turbulence, and ii) assume that the jump conditions for isothermal MHD shocks determine the scaling of the typical size of protostellar cores. Based on these fundamental assumptions, we are  able to derive a  power law mass distribution of gravitationally unstable  cores, with a slope consistent with the Salpeter stellar IMF (Salpeter  1955). In  order  to  derive  the  mass  distribution  of gravitationally unstable cores  below one solar mass we  also make use of  the Probability  Density Function  (PDF)  of the  mass density  in super--sonic turbulence, following Padoan, Nordlund \\& Jones (1997).  The  formation of dense  cores and  the relation  between the  size of cores and the thickness of the postshock gas are discussed in the next section.  In  \\S~3 the power spectrum  of super--sonic turbulence is discussed.  Results on  the PDF of the gas density in super--sonic turbulence are summarized in \\S~4.  The mass distribution of dense  cores formed  by the process  of turbulent  fragmentation is derived in \\S~5, and the mass distribution of gravitationally unstable cores is  computed in  \\S~6.  The relevance  of these results  for the origin of the stellar IMF is  discussed in \\S~7 and 8. Conclusions are summarized in \\S~9. ",
        "conclusions": "In conclusion, we have related the stellar IMF to the mass distribution of dense cores formed by the process of turbulent fragmentation, assuming that only cores with gravitational energy in excess of their magnetic and thermal energy can collapse as protostars. Most sub--critical cores disperse back into the turbulent flow, and are therefore irrelevant for the process of star formation, as we have argued in other recent works (Padoan et al.\\ 2001a, b). Previous theories of star formation (see Shu, Adams \\& Lizano 1987) assume that stars of small and intermediate mass are formed from sub--critical cores. In such theories, sub--critical cores are in static equilibrium and evolve quasi--statically, on the time--scale of ambipolar drift. We have argued that such a scenario is inconsistent with the turbulent nature of MCs (Padoan et al.\\ 2001a, b).  We have derived the mass distribution of dense cores generated by the process of turbulent fragmentation, and have found a power law mass distribution consistent with the Salpeter stellar IMF,  for stellar masses larger than  1--2 m$_{\\odot}$. This result is a direct consequence of fundamental physical properties of super--sonic turbulence in MCs, such as the power spectrum of turbulence and the jump conditions for isothermal MHD shocks. We have also  shown that another fundamental  physical property of turbulent flows, namely the PDF of the gas density, can explain the shape of the IMF at smaller stellar masses, and the formation of gravitationally unstable cores of sub--stellar mass.  The main results of this work are: i) The power law stellar IMF at masses larger than 1--2 m$_{\\odot}$ is the result of the near self--similar nature of inertial range super--sonic turbulence; ii) The low mass roll-over and cut-off of the IMF is caused by the combined effects of the thermal support of the smallest cores against gravitational collapse and the density PDF of super--sonic turbulence; iii) The mass--scaling of the peak of the IMF may be expressed as a function of the physical parameters of turbulent star--forming clouds; primarily their rms turbulent velocity, temperature, magnetic field strength and average density; iv) For physical  parameters typical of nearby MCs  the peak  of the  IMF is predicted  to be  around 0.3--0.5 m$_{\\odot}$;  v) The slope  of the  IMF at masses  larger than  1--2 m$_{\\odot}$  is determined by the  inertial range  spectral  index of super--sonic  turbulence and  the jump  conditions for  isothermal MHD shocks;  vi) For  magnetically dominated jump conditions, which  are applicable  to typical  molecular cloud  conditions, the  IMF spectral index is equal  to $3/(4-\\beta)$, where $\\beta$ is  the inertial range spectral index  of super--sonic turbulence; vii) For a value of the inertial range spectral index $\\beta=1.74$, consistent with Larson's relations and with new numerical results, the IMF spectral index is $x=1.33$, almost identical to Salpeter's slope.  We conclude that the process of turbulent fragmentation is essential to the origin of the stellar IMF, in support of the thesis that star formation can be viewed as the main consequence of the dissipation of super--sonic turbulence in molecular clouds."
    },
    "0011/astro-ph0011379_arXiv.txt": {
        "abstract": "The past, present and future of cosmic microwave background (CMB) anisotropy research is discussed, with emphasis on the Boomerang and Maxima balloon experiments. These data are combined with large scale structure (LSS) information derived from local cluster abundances and galaxy clustering  and high redshift supernova (SN1)  observations to explore the inflation-based cosmic structure formation paradigm. Here we primarily focus on a simplified inflation parameter set, $\\{\\omega_b,\\omega_{cdm},\\Omega_{tot}, \\Omega_Q,w_Q,n_s,\\tau_C, \\sigma_8\\}$. After marginalizing over the other cosmic and experimental  variables, we find the current CMB+LSS+SN1 data gives $\\Omega_{tot} = 1.04\\pm 0.05$, consistent with (non-baroque) inflation theory. Restricting to $\\Omega_{tot}=1$, we find a nearly scale invariant spectrum, $n_s =1.03 \\pm 0.07$. The CDM density, $\\omega_{cdm}=0.17 \\pm 0.02$, is in the expected range, but the baryon density, $\\omega_b\\equiv \\Omega_b {\\rm h}^2 = 0.030\\pm 0.004$, is slightly larger than the current $0.019\\pm 0.002$ Big Bang Nucleosynthesis estimate. Substantial dark (unclustered) energy is inferred, $\\Omega_Q \\approx 0.68 \\pm 0.05$, and CMB+LSS $\\Omega_Q$ values are compatible with the independent SN1 estimates. The dark energy equation of state, parameterized by a quintessence-field pressure-to-density ratio $w_Q$, is not well determined by CMB+LSS ($w_Q < -0.3$ at 95\\% CL), but when combined with SN1 the resulting $w_Q < -0.7$ limit is quite consistent with the $w_Q$=$-1$ cosmological constant case. Though forecasts of statistical errors on parameters for current and future experiments are rosy, rooting out systematic errors will define the true progress. \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011523_arXiv.txt": {
        "abstract": "We investigate whether gamma-ray pulsars viewed at a large angle to the neutron star magnetic pole could contribute to the new population of galactic unidentified EGRET sources associated with the Gould Belt. The faint, soft nature of these sources is distinctly different from both the properties of unidentified EGRET sources along the galactic plane and of the known gamma-ray pulsars.  We explore the possibility, within the polar cap model, that some of these sources are emission from pulsars seen at lines of sight that miss both the bright gamma-ray cone beams and the radio beam.  The off-beam gamma-rays come from high-altitude curvature emission of primary particles, are radiated over a large solid angle and have a much softer spectrum than that of the main beams.  We estimate that the detectability of such off-beam emission is about a factor of 4-5 higher than that of the on-beam emission.  At least some of the radio-quiet Gould Belt sources detected by EGRET could therefore be such off-beam gamma-ray pulsars. GLAST should be able to detect pulsations in most of these sources. ",
        "introduction": "The origin of the 170 unidentified gamma-ray point sources detected by the Energetic Gamma-Ray Experiment Telescope (EGRET) on board the Compton Gamma-Ray Observatory (CGRO) throughout its mission (Hartman et al. 1999) is one of the most intriguing questions of high-energy astrophysics.  Grenier \\& Perrot (1999) found that some of the unidentified sources are significantly correlated with the Gould Belt of massive stars, a nearby galactic structure surrounding the Sun.  The Gould Belt is an expanding disk of gas and young stars, most with ages less than 30 million years, inclined about $20^0$ to the galactic plane. Recently, Gehrels et al. (2000) studied the spatial and flux distribution of the 120 steady sources in the third EGRET catalog and found that they divide into two groups: higher flux sources distributed along the galactic plane with latitudes less than $5^0$, and a new population of lower flux sources at mid-latitude which indeed seem to be associated with the Gould Belt, at a relatively nearby distance of 100-300 pc.  The two groups also show significant differences in their spectra, with the low-latitude sources having average photon spectral index of $2.18 \\pm 0.04$ and the mid-latitude sources having average photon index of $2.49 \\pm 0.04$. None of these sources have known counterparts at other wavelengths. Grenier \\& Perrot (1999) suggested that the Gould Belt $\\gamma$-ray sources could be young pulsars, formed in supernova explosions of massive stars of the belt. However, the relatively soft spectra and luminosities of these sources, averaging $\\sim 6 \\times 10^{30}\\,{\\rm erg\\,s^{-1}}$, are not characteristic of the known $\\gamma$-ray pulsars detected by EGRET, which have spectral indices of $1.5 - 2.0$, luminosities $10^{32} - 10^{34}\\,{\\rm erg\\,s^{-1}}$ and 0.5-4 kpc distances.  Here, we show that the bulk of the Gould-Belt $\\gamma$-ray sources could be pulsars if their emission is seen at large angles to their magnetic axes, such that we are missing the bright, hard $\\gamma$-ray beams but detecting only the off-beam emission. According to the polar cap model (e.g. Daugherty \\& Harding 1996 [DH96]), $\\gamma$-ray emission occurs throughout the entire pulse phase. A synchrotron-pair cascade at low altitude radiates the hard on-beam emission in a hollow cone centered on the magnetic pole to produce the bright double-peaked pulses.  Primary electrons continue to radiate curvature emission on open field lines to high altitudes beyond the cascade region, producing a lower level of softer off-beam emission.  Due to the flaring of the dipole field lines, this emission may be seen over a large solid angle, far exceeding that of the main beams.  Since the radio emission is expected to originate within ten stellar radii of the neutron star surface, it is quite probable to see off-beam $\\gamma$-ray emission and miss the radio beam.  Such off-beam or off-pulse emission may have been detected by EGRET in the Crab, Vela and Geminga $\\gamma$-ray pulsars during non-pulse phases (Fierro et al. 1998).  We show that this expected off-beam emission has the right characteristics (spectral index, luminosity and numbers) to account for the EGRET sources associated with the Gould Belt. ",
        "conclusions": "With the polar cap cascade geometry inferred from the known gamma-ray pulsars, we find that there is a large phase space for which the line-of-sight misses both the main gamma-ray beam and possibly the radio beam, but still cuts across a much broader, fainter gamma-ray beam produced by the curvature-radiation cooling of the primary particles.  We identify such off-beam $\\gamma$-ray pulsars as candidates for the new population of the unidentified EGRET sources associated with the Gould Belt.  The off-beam polar cap cascade emission exhibits the low luminosity and soft spectrum in the EGRET band that is characteristic of the observed Gould Belt sources. A rough estimate indicates that the off-beam sources might be $\\sim 4 - 5$ times more detectable than the on-beam sources due to the larger solid angle of the off-beam emission. A more accurate prediction of the fraction of off-beam $\\gamma$-ray pulsars in the Gould Belt will require modeling of the pulsar population and luminosity selection effects. However, since EGRET detects all favorably oriented on-beam pulsars and a large fraction of off-beam pulsars in the belt, luminosity selection effects may not be a dominant factor. GLAST will be able to detect pulsed emission from these sources, and the pulse shape should be a broad single-peaked profile.  Off-beam emission is not expected in outer gap models (Yadigaroglu \\& Romani 1995, Cheng \\& Zhang 1998), which predict that most low-latitude ($|b|<5^{\\rm o}$) unidentified EGRET sources in the plane are radio-quiet, Geminga-like pulsars.  Thus, the radio-quiet outer-gap pulsars will be on-average as luminous in $\\gamma$-rays as the radio-loud pulsars.  This is in contrast to our prediction that polar cap radio-quiet pulsars (at least those detected within a few hundred pc) should be on average less luminous in $\\gamma$-rays than radio-loud pulsars.  We thank Isabelle Grenier, Matthew Baring and Seth Digel for valuable input and discussions."
    },
    "0011/astro-ph0011009_arXiv.txt": {
        "abstract": "The observed quasi-periodic oscillations (QPOs) in low-mass X-ray binaries have been suggested to be probably related to the characteristic % oscillation modes in the inner part of a relativistic accretion disk. In this paper we study the trapped feature of these oscillation modes by considering the radial perturbations and viscosity effects. We find that the trapped oscillations are amplified significantly when the viscosity parameter increases and the trapped size varies when the radial wavenumber and Mach number change. Our results support that the stable hundred-hertz QPOs and their low fractional rms amplitudes found in two Galactic microquasars may be explained by some trapped non-damping oscillation modes in the inner part of accretion disks around  rotating black holes. ",
        "introduction": "In recent years, hundred-hertz QPOs % have been observed  in several black hole X-ray binaries with {\\it RXTE} satellite. These high-frequency QPOs are probably similar to the kilohertz QPOs  observed frequently in a dozen  neutron star low-mass X-ray binaries. However,  the high frequency QPOs in two Galactic microquasars,  GRO J1655-40 and GRS 1915+105,  are very stable in frequency at 300 Hz and 67 Hz respectively (Morgan, Remillard \\& Greiner 1997; Remillard \\etal 1999). These stable QPOs have also very low fractional rms amplitude ($\\sim 1 \\%$). Theoretical models involving frame-dragging (Cui, Zhang \\& Chen 1998), diskoseismic g-mode (Nowak \\etal 1997), and periastro precession (Stella, Vietri \\& Morsink 1999) have been proposed to explain these stable QPOs.  But the origin of them is still uncertain.  Recently Psaltis \\& Norman (2000; hearafter PN) suggested that the observed QPOs in X-ray binaries may be related to the strong resonances peaked at some characteristic oscillation frequencies in the inner part of accretion disks. However, they introduced a simplified term for angular momentum transfer and  did not % discuss  % the damping and growth rates of oscillation modes. % It is also important to check if some of these modes can be trapped  in the inner  disk and  if the low fractional rms amplitude of QPOs observed in two microquasars can be explained by the trapped feature of oscillation modes. ",
        "conclusions": "{\\it RXTE} observations on GRO J1655-40 have shown 300 Hz QPOs with $\\sim 1\\%$ rms magnitude. If this stable QPO originates from the trapped disk oscillations, from Figures 1 and 2 we see that it may be explained by either the lower-frequency acoustic mode in the case of $m=1$ and $n=0$ or the lower-frequency corrugation mode in the case of $m=n=1$ and $a_*\\sim 0.9$. However, the rapidly spinning black hole is more favored to explain the high QPO frequency. The unambiguous determinations of the inner disk radius and the black hole spin are still needed to confirm this explanation. Further studies on the origin of  low-frequency QPOs and their correlation with high-frequency QPOs, as well as the correlation between  spectral and temporal properties, will  help us to understand more clearly the origin  of stable QPOs in microquasars."
    },
    "0011/astro-ph0011264_arXiv.txt": {
        "abstract": "We present a self-similar, steady-state model describing both a magnetized accretion disc and a magnetohydrodynamic jet. We focus on the role of a hot corona in such a structure. This corona enables the disc to launch various types of jets. By considering the energy conservation, we also present a diagnostic of the luminosity of the magnetized disc, which could explain some observational signatures of galactic objects. ",
        "introduction": "In the Universe, jets of matter are observed around almost every kind of systems where plasma accretes onto a central object. The central object can be a young protostar, a compact galactic object or a supermassive black hole. The jet velocity is always found to be a few times the escaping velocity of the central object (e.g., \\opencite{livio}). The most successful model, accounting for all known properties of observed jets, relies on a large scale magnetic field anchored on an accretion disc (\\opencite{bland}). The field extracts the disc angular momentum, thereby allowing accretion, and accelerates a fraction of disc material to high speeds. Moreover, the magnetic field confines the accelerating jet matter, leading to a cylindrical collimation of the jet.  Hereafter, such systems are called Magnetized Accretion-Ejection Structures (MAES).  In a MAES, the bipolar magnetic field pinches {\\it the disc}. Thus, the only force able to counteract both the magnetic and gravitational forces, {\\it whithin the disc},  is the vertical plasma pressure gradient. As a consequence, the mass loading in the jet depends critically on this vertical equilibrium. This equilibrium is the connection between the accretion motion and the jet. A more detailed explaination of this phenomenon is displayed in \\inlinecite{fab3}. ",
        "conclusions": ""
    },
    "0011/astro-ph0011052_arXiv.txt": {
        "abstract": "The strong activity of radio galaxies should have led to a nearly ubiquitous presence of fossil radio plasma in the denser regions of the inter-galactic medium as clusters, groups and filaments of galaxies. This fossil radio plasma can contain large quantities of relativistic particles (electrons and possibly protons) by magnetic confinement.  These particles might be released and/or re-energized under environmental influences as turbulence and shock waves. Possible connections of such processes to the formation of the observed sources of diffuse radio emission in clusters of galaxies (the cluster radio halos and the cluster radio relics) are discussed. ",
        "introduction": "The remnant of the radio lobes of a former active radio galaxy get rapidly invisible for radio telescopes, justifying the name {\\it 'radio ghosts'} for such patches of fossil radio plasma (\\cite{1999dtrp.conf..275E}). Radio ghosts should be an important ingredient of the inter galactic medium (IGM). The present space density of one or a few radio galaxies per galaxy cluster was exceeded by orders of magnitude during the epoch of violent quasar activity around $z = 2$. There are several routes to estimate the energy released by earlier radio galaxies. One is to integrate the evolution of the radio-luminosity-function, translated to jet-power via an empirical jet-power--radio-luminosity relation, based on equipartition energy densities of radio lobes. The result is (\\cite{2000A&A...360..417E}) \\begin{equation} E_{\\rm jet} = 3\\cdot 10^{66} \\,{\\rm erg\\,Gpc^{-3}} \\,(f_{\\rm eq}/3)\\,, \\end{equation} where $f_{\\rm eq}$ gives the ratio of true to equipartition energy density in the radio plasma, which is likely bigger than one. This amount of energy corresponds to $\\sim 0.1$ keV per cosmic baryon, and is therefore not negligible, since large-scale structure formation has produced $\\sim 1$ keV/baryon by gravitational infall.  An independent route is to use the Magorrian relation ($\\eta_{\\rm bh} = M_{\\rm bh}/M_{\\rm bulge} = 0.002$) between the galaxy bulge ($M_{\\rm bulge}$) and central black hole ($M_{\\rm bh}$) mass. Assuming an efficiency of $\\varepsilon_{\\rm jet} = 0.1$ of rest mass to jet energy conversion of the accretion process feeding the black hole growth, one derives from the observed present day galaxy population (\\cite{1998AA...333L..47E}) \\begin{equation} E_{\\rm jet} = 3\\cdot 10^{67} \\,{\\rm erg\\,Gpc^{-3}} \\,(\\eta_{\\rm bh}/0.002)\\,(\\varepsilon_{\\rm jet} /0.1)\\,. \\end{equation} If this estimate is correct, this amount of energy could easily provide the cosmic entropy-floor of 1 keV/baryon, which seems to be required in order to explain the X-ray luminosity--temperature relation of groups and clusters of galaxies.  Finally, a third way to estimate the energy in radio plasma is to assume that the efficiencies of accretion discs in active galactic nuclei (AGN) to produce X-rays and radio jets are the same. Assuming further that the observed X-ray background results only from unresolved AGN, one gets (\\cite{1998AA...333L..47E}) \\begin{equation} E_{\\rm jet} = 3\\cdot 10^{67} \\,{\\rm erg\\,Gpc^{-3}} \\,(\\varepsilon_{\\rm jet} /\\varepsilon_{\\rm X-ray})\\,. \\end{equation}   This should illustrate the importance of fossil radio plasma for IGM properties. A discussion of the possible influences of radio ghosts on various aspects of extra-galactic astrophysics can be found in \\cite{1999dtrp.conf..275E} and \\cite{Pune99}. Here, recent progress on two of these aspects should be briefly reported: an attempt to estimate the escape rate of relativistic particles from radio ghosts (\\cite{crdiff2000}), and a new model for the revival of of fossil radio plasma (\\cite{ensslin2000b}). ",
        "conclusions": ""
    },
    "0011/astro-ph0011578_arXiv.txt": {
        "abstract": "The 10$^{51}$ ergs released in a supernova have far reaching consequences in the galaxy, determining elemental abundances, accelerating cosmic rays, and affecting the makeup of the interstellar medium. Recently the spectra of several supernova remnants have been found to be dominated by nonthermal emission. Separating the thermal and nonthermal components is important not only for the understanding of cosmic-ray acceleration and shock microphysics properties but for accurate assessment of the temperatures and line strengths. New models designed to model spatially resolved synchrotron X-rays from type Ia supernovae can contribute to the understanding of both the thermal physics (dynamics, abundances) and nonthermal physics (shock acceleration, magnetic-field amplification) of supernova remnants. I will describe model fits to SN 1006, emphasizing the physical constraints that can be placed on SNRs, abundances, and the cosmic-ray acceleration process. ",
        "introduction": "\\begin{figure} \\plotfiddle{kdyer_1.ps}{2.5in}{0}{50}{50}{-150}{-170} \\caption{The radio flux measurement with a power law spectrum extrapolated to X-ray energies. The X-rays are the unfolded ASCA spectrum, shown for comparison. \\label{fluxes}} \\end{figure}  A growing number of SNRs (listed in Table 1) are thought to show evidence for synchrotron X-rays, for a variety of reasons. Most are suspected of nonthermal X-rays because they show hard energy excesses that can be fit with powerlaws. A few  (SN~1006, G347.3-0.5) have nearly featureless X-ray spectra which are clearly dominated by nonthermal emission. Some remnants, while not dominated by nonthermal emission in the ASCA band, show other evidence pointing to its presence. OSSE observations of Cassiopeia A at 400-1250 keV (The 1996) are well described by a broken power law, steepening to higher energies (Allen 1997). RCW 86 is an especially interesting case. It shows anomalously low abundances when fit with thermal models (Vink 1997), but the addition of a synchrotron model bring the abundances to within expected ranges (Borkowski 2001).       The most straightforward explanation for nonthermal X-ray emission (non-plerionic in these SNRs) is synchrotron emission from the same mechanism that is responsible for the radio emission. It is straightforward to determine the synchrotron spectrum from the electron energy population However, the extrapolation from the radio spectrum measured by single dish instruments, over-predicts the X-ray flux (see Figure 1) -- there must be a mechanism that causes the electron population to roll off at the highest energies. Reynolds \\& Chevalier (1981) attributed this rolloff to synchrotron losses.   The sharpest physically plausible spectral rolloff would be the emission resulting from an exponential cutoff of the electron power-law energy distribution, radiating in a uniform magnetic field. This was embodied in the XSPEC model {\\it SRCUT} and used by Reynolds \\& Keohane (1999) to set limits on the energies to which SNRs could accelerate electrons (with implications for cosmic ray ions). A more sophisticated treatment requires addressing the mechanism limiting electrons at the highest energies. The cutoff could be due to the age of the SNR, radiative losses from synchrotron or inverse-Compton emission, or electron escape(Reynolds 1998).     \\begin{table} \\caption{SNRs which have reported nonthermal X-rays\\label{table1}} \\begin{tabular}{lcc} \\tableline SNR & Instrument & Nonthermal reference \\\\ \\tableline Cassiopeia A\t& OSSE, RXTE, BeppoSAX & The 1996, Allen 1997, Favata 1997\\\\ RCW 86\t\t& ASCA\t\t\t\t& Borkowski 2001, Bamba 2001\\\\ SN 1006\t\t& ASCA\t\t\t& Koyama 1995, Dyer 2001\\\\ Kepler \t\t& ASCA, XTE\t\t\t& Decourchelle 1999\\\\ G347.3-0.5\t& ASCA \t\t\t\t& Koyama 1997, Slane 1999\\\\ G327.1-1.1\t& ASCA, ROSAT\t\t\t& Sun 1999\\\\ \\tableline \\tableline \\end{tabular} \\end{table} ",
        "conclusions": "We have found, both in integrated and spatially resolved models, that the measured abundances and temperature change with the addition of accurate synchrotron models. In the future, accurate accounting for a possible nonthermal component will be a necessary precondition to using nucleosynthesis model predictions to interpret abundances in any SNR. {\\it SRCUT} and {\\it SRESC}, new models in XSPEC 11, accurately describe this emission and are a substantial improvement over powerlaw models both in terms of physical description and fitting.     All regions of SN 1006, even the polar caps which show obvious thermal lines, seem to have significant nonthermal emission. We also find that the spectra of the north and the south limb are not identical. While most of the emission from the center does seem to be thermal, the fit can tolerate a significant amount of nonthermal emission.  Technology is surpassing theory -- Chandra arcsecond-resolution observations of SNRs including SN~1006 will soon be available, requiring models capable of describing extremely small scale variations. These synchrotron models will be useful both for analyzing higher resolution observations of SN~1006, and for interpreting other SNRs with nonthermal X-ray emission.   \\begin{figure} \\plotfiddle{kdyer_4.ps}{237pt}{-90}{50}{50}{-190}{275} \\caption{A limb version of the SRESC model fit to ASCA data from the south limb of SN~1006. Note the line-like residuals. \\label{needthermal}} \\end{figure}  \\begin{figure} \\plotfiddle{kdyer_5.ps}{237pt}{-90}{50}{50}{-190}{275} \\caption{The limb version of the SRESC model fit to data from the south limb of SN~1006 with a plane shock (VPSHOCK, Borkowski 2001). \\label{thermallimb}} \\end{figure}"
    },
    "0011/astro-ph0011322_arXiv.txt": {
        "abstract": "The evolution of Ly$\\alpha$ absorber counts is simulated for a model population of absorbing galaxies. The distribution of gas relative to galaxies could evolve between moderate and low redshifts, but constraints are needed on the strength and evolution of the ionizing UV background. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011114_arXiv.txt": {
        "abstract": "indent} \\def\\titlehang{\\hangindent\\abstractindent \\hangafter 0 \\relax}  \\newdimen\\secskipamount  \\secskipamount=1pt \\newdimen\\subskipamount  \\subskipamount=1pt   \\newdimen\\bottomtol \\bottomtol=0.03\\vsize \\def\\secskip{\\par \\ifdim\\lastskip<\\secskipamount \\removelastskip \\fi \\vskip 0pt plus \\bottomtol \\penalty-250 \\vskip 0pt plus -\\bottomtol \\relax \\vskip\\secskipamount plus3pt minus3pt} \\def\\subskip{\\par \\ifdim\\lastskip<\\subskipamount \\removelastskip \\fi \\vskip 0pt plus 0.5\\bottomtol \\penalty-150 \\vskip 0pt plus -0.5\\bottomtol \\relax \\vskip\\subskipamount plus2pt minus2pt} \\def\\subsubskip{\\par \\ifdim\\lastskip<\\subskipamount \\removelastskip \\fi \\vskip 0pt plus 0.5\\bottomtol \\penalty-150 \\vskip 0pt plus -0.5\\bottomtol \\relax \\vskip\\subskipamount plus2pt minus2pt \\hskip 10pt} \\long\\def\\aaabstract#1{\\centerline{\\null} \\vskip 1.52cm {\\absnarrower \\noindent {\\bf Summary.} #1 \\par} \\oneskip \\oneskip}   \\outer\\def\\unnumberedsectionbegin #1 #2\\par {\\secskip \\noindent {{\\bf  #1} \\dotfill #2} \\nobreak \\vskip 1pt \\noindent}  \\outer\\def We have observed a sample of OH 1612--MHz masing objects in all four OH ground--state transitions with the Australia Telescope Compact Array. One likely post--AGB object is found to emit in the \\sat , \\mla\\ and \\sath\\ transitions. We discuss the evidence that this object may be an early post--AGB object and the possibility for such a circumstellar envelope to harbour a \\sath\\ maser. We argue that during a very brief period, just after the star has left the thermally--pulsing phase of the AGB and the wind velocity starts to increase, post--AGB objects might show \\sath\\ emission. The best objects to search for such emission would be those that are masing at 1612 MHz and 1665 MHz, but not at 1667 MHz nor in the 22--GHz H$_2$O transition. ",
        "introduction": "OH maser emission from the four ground-state transitions at 1612.231, 1665.402, 1667.359 and 1720.530 MHz occurs in a wide range of objects. Extra--galactic OH--maser emission is associated with active galactic nuclei and enhanced central star formation (Baan \\& Haschick 1987; Randell \\etal 1995), with strongest emission, in general, from the main lines at 1665/1667 MHz. OH maser emisson from galactic star forming regions (SFRs) is also strongest in the main lines, with dominant emission mostly at 1665 MHz.  The satellite line at 1720 MHz is best known for its occurence in supernova remnants, whereas that at 1612 MHz is strongest in the optically--thick envelopes of oxygen--rich asymptotic--giant--branch (AGB) stars. Mainline OH emission may be present in optically--thin envelopes; from nearby Mira variables, with low mass-loss rates around 10$^{-7}$ M$_\\odot$ yr$^{-1}$, only mainline emission is detected.  At the end of the AGB evolutionary phase, the star undergoes a transformation which takes a couple of thousand years, during which it evolves to become a planetary nebula (PN). During this post-AGB stage, the mass-loss rate decreases by a factor of 1000 or more whilst the stellar surface temperature rises by a factor of 10 or 20 to above 50,000 K.  The star changes from losing mass in a cool, dense wind to losing mass in a much hotter, lower density wind.  Strong changes also occur in the envelope kinematics and morphologies. In particular, whilst the circumstellar envelopes of AGB stars are generally spherically symmetric, images of planetary nebulae show axi-symmetric morphologies in over 50\\% of cases (Manchado \\etal2000).  During the initial searches in the 1970s it was quickly found that \\sath\\ emission was not detected from AGB stars. As this agreed with theoretical expectations (eg.~Elitzur, Goldreich \\& Scoville 1976), there have been no systematic searches for this maser in stellar sources. However, it is imaginable that the shocks in some post--AGB stars could temporarily induce \\sath\\ maser emission. We therefore decided to observe, for the first time, a modest sample of \\sat --selected sources, including post--AGB objects, at 1720 MHz. A total of 18 objects were selected from Sevenster \\etal(1997a,b) on the basis of their unusual OH 1612--MHz profiles. Observations in all four ground--state OH transitions were taken with the Australia Telescope Compact Array (ATCA) in November 1998 (1612, 1665 and 1667 MHz) and September 1999 (1720 MHz). Eleven of the sources are likely post--AGB objects.  The full results from this study will be published separately. Here we report the detection of 1720 MHz from the source \\Btnt\\ (IRAS 18043$-$2116). In \\Sct 2, we present the observations and in \\Sct 3 the evidence that \\Btnt\\ is a post--AGB object. In \\Sct 4 we discuss a possible scenario for the \\sath\\ excitation. Conclusions are given in \\Sct 5. ",
        "conclusions": "We have detected an OH \\sath\\ maser in a likely post--AGB envelope. From theoretical as well as observational arguments, it is possible that this maser exists briefly in such environments, at least in some sources. We argue that the \\sath\\ maser could arise in the region where the AGB-- and the starting PN winds collide. We expect that this is a short-lived event and that our detection rate of 1/11 is due to favourable small number statistics. The most likely objects to show this emission may be those that have fairly strong emission at 1612 MHz, with weaker \\mla\\ and possibly SiO maser emission, but no \\mlb\\ and certainly no H$_2$O 22--GHz maser emission."
    },
    "0011/astro-ph0011358_arXiv.txt": {
        "abstract": "{We determine that the sizes of bursts in mean-square current density in a reduced magnetohydrodynamic (RMHD) simulation follow a power-law probability density function (PDF). The PDFs for burst durations and waiting time between bursts are clearly not exponential and could also be power-law. This suffices to distinguish their behaviour from the original Bak et al.\\ sandpile model which had exponential waiting time PDFs. However, it is not sufficient to distinguish between turbulence, some other SOC-like models, and other red noise sources.} ",
        "introduction": "The widespread occurrence of both self-affine time series with ``$1/f$'' power spectra and spatial fractals in nature led \\cite{BakEA87,BakEA88} (BTW) to propose the hypothesis of self-organised criticality (SOC) \\citep{BakSOC,JensenSOC,SornetteSOC} as their common origin. Their proposal was based on the demonstration of a  ``sandpile'' cellular automaton (see also the earlier work of \\cite{Katz86}) which appeared to be attracted from arbitrary initial conditions (``self-organisation'') to a critical state characterised by fluctuations on all scales in the energy released by the system (``criticality''). Power-law probability density functions (PDFs) for the sizes and durations of energy bursts were the main observed signatures of criticality, and were tested by finite-size scaling of the PDFs with system size \\citep{CardySR}.  One of the original applications proposed by BTW for their idea was fully developed turbulence, in view of the scaling behaviour of such systems, and the intermittency of their energy dissipation. Furthermore, intermittent turbulence has been an inspiration for later ``sandpile''-like cellular automata such as the forest fire model of \\cite{BakEA90}. The SOC paradigm has since found many applications \\citep{BakSOC,JensenSOC,SornetteSOC}, one of which has been its use by \\cite{LuHamilton91} and subsequent authors to explain the observed power-law distributions for the magnitudes, intensities, and durations of solar flares. SOC has since also been applied to other natural and artificial plasma confinement systems, notably the Earth's magnetosphere (a recent review is that of \\cite{ChapWat01}) and tokamaks (e.g. \\citep{ChapEA01}).   In turn, recent studies in the solar flare context, of both shell models \\citep{BoffettaEA99} and simulations based more directly on the magnetohydrodynamic (MHD) equations \\citep{GeorgoulisEA98,EinaudiVelli99} have focused attention on the fundamental questions of the ways in which SOC and turbulence may differ \\citep{BoffettaEA99} or, conversely, the extent to which SOC may serve as a model for turbulence \\citep{EinaudiVelli99}. These more general questions are our main focus in this paper.  Also included is some brief discussion of the application of these ideas to solar flares and plasma turbulence in the solar wind, magnetosphere, and elsewhere.   Recent work has shown that magnetically forced 2D MHD turbulence produces power-law PDFs in the size and duration of bursts in  spatially averaged Ohmic energy dissipation $\\langle \\eta J^2 \\rangle$ \\citep{GeorgoulisEA98,EinaudiVelli99} (G98). Recall that such power-law PDFs are a necessary but not sufficient condition for SOC \\citep{JensenSOC}. In addition \\cite{EinaudiVelli99} showed that a cellular automaton with rules chosen to be consistent with the MHD model also produces avalanches (and power-law PDFs).   In contrast, \\cite{BoffettaEA99} and \\cite{GiulianiEA00} (B99) attributed the presence of a power-law PDF in the observed waiting time between solar flares  to turbulence. Rather than simply being due to  the scale-free turbulent cascade itself, B99's suggested mechanism for the production of power law waiting time PDFs was on-off intermittency. They noted that simple prototype models of on-off intermittency and a many-oscillator shell model of turbulence both had waiting time PDFs which were power-law, while the original BTW sandpile algorithm did not, having instead an exponential PDF of waiting times.  Two important criticisms of B99's model and its interpretation have recently been made in the context of solar flares. \\cite{EinaudiVelli99} pointed out that its attractors were states in which velocity and magnetic field were aligned, whereas they asserted that force free (aligned current and magnetic field) states were probably more appropriate to the corona. \\cite{Wheat00} showed how a Poisson process varying on a long time scale (e.g. the quasi-periodic solar cycle) could convert an exponential waiting time distribution such as that from the original sandpile model into a power-law one of the type observed.  These results are however  less relevant to  the more general plasma turbulence case. In physical systems where there is no long-term periodicity to motivate the ``periodic-Poisson\" assumption of \\cite{Wheat00} there is no apparent reason to prefer such a mechanism for observed power laws to an intrinsically scale-free one. In addition, the work of \\cite{Wheat00} seems  to have been partly been a response to B99's over-general assertion that all SOC models \\emph{must} have exponential waiting time PDFs. This is not true in general (c.f. the discussions in \\cite{Galtier01} and \\cite{FreemanEA00b}; and the models studied by  \\cite{PaczuskiEA96}). In consequence, several interesting questions remain open.    One is whether the burst size and duration PDFs found in simulations of 2D MHD turbulence \\citep{GeorgoulisEA98,EinaudiVelli99} are also observed for either full MHD or reduced MHD (RMHD), in which ``slow'' $k_z$ dependence is retained, (see below).  The second is whether the power-law waiting time PDFs seen in both the B99 one-dimensional shell model and the \\cite{Galtier01} 1D MHD simulation, are also seen in higher dimensions.  Finally there is the question of whether a minimal set of scale-free ``burst'' PDFs and power spectra can be identified which suffices to identify SOC (or turbulence) in a physical system. This last question is also highly topical in magnetospheric and laboratory plasmas \\citep[cf.\\ ][]{Krommes00,FreemanEA00b,KovacsEA01}.  In this paper we address these questions by examining time series of various spatial averages of the squared electric current density, $j^2(x,y,z,t)$, drawn from a RMHD simulation. A threshold method, used by previous authors to construct burst-size PDFs, is applied to the time series. Power-laws in size (and arguably also in duration) are found, extending the forced 2D MHD results of \\cite{GeorgoulisEA98} and \\cite{EinaudiVelli99} to forced RMHD. The PDF of waiting times is also not exponential, confirming that higher-dimensional RMHD is in keeping with the predictions of B99 based on a 1D shell model. Because the fixed threshold method employed detects fractality, which was a predicted feature of SOC but is also generic to red noise, we then consider to what extent the current evidence is unambiguous. We conclude by suggesting a direction for future research. ",
        "conclusions": "We find that the PDF of burst size measured in our simulation is  power-law in form, independent of the choice of threshold. The PDFs of duration and waiting time appear to have the same basic form (although the evidence is much less clear cut). Intriguingly, we have observed this property  in a time series (Figure 1a) whose PDF (Figure 1b)  is not long-tailed but rather symmetric. Power-law burst PDFs are often seen coupled with long-tailed underlying PDFs, and so it is sometimes thought that the scaling range of a signal will be controlled by the mean to peak ratio (governed by $\\sigma$ and $\\mu$ for Gaussian data, $\\mu$ for Poisson data, etc.)  In fact, however, burst size is also governed by the degree of persistence $\\beta$ in the signal, because it is not just affected by the probability of N points being above a line but is controlled by the probability of N non-independent, successive points all being above the line. \\cite{MalTurc99} have noted that $\\beta$ can be varied independently of the PDF of the members of a time series, so that even a time series with a Gaussian distribution of amplitudes but a non-zero $\\beta$ would have a non-negligible probability of several successive values exceeding a threshold. In addition, for such a series, the scaling properties of the distributions of waiting time and inter-burst interval are set entirely by $\\beta$ and would be power laws.  This raises an interesting question, however  (see also \\cite{FreemanEA00b}). Red noise (``$1/f$'') time series of the type which SOC systems were originally expected to produce, and whose appearance SOC was proposed to account for, are fractal. The distribution of isosets\\footnote{Defined as the set of times at which the time series crosses a fixed level.} of such a time series is a power-law \\citep{AddisonFractals}. Hence the burst duration and waiting time PDFs of red noise when found by a threshold method should be power-laws, regardless of whether the noise is produced by turbulence or one of the class of SOC-type models which do produce red noise. The original BTW model can be eliminated as its waiting time series was later shown to be uncorrelated and thus not red noise \\citep{JensenSOC,BoffettaEA99,FreemanEA00b}. To determine if a process is SOC in the sense of BTW's original proposal, one needs information about spatial correlations as well as temporal correlations.  This is because SOC was proposed as a mechanism linking spatial fractality and temporal persistence (``1/f\" noise).  This reinforces the point that, rather than simply temporal information (burst durations and waiting times) or avalanche distributions (``burst sizes\"), to differentiate between turbulence and SOC it will be necessary to make unambiguous predictions about spatial structure for each phenomenon, requiring at minimum then availability of spatial correlation functions. We note that correlation between bursts e.g. the debated ``sympathetic\" quality of solar flares \\citep{WheatEA98} could in principle occur through time or space correlation.  This line of investigation will be pursued in future papers."
    },
    "0011/astro-ph0011391_arXiv.txt": {
        "abstract": "s{ Rich superclusters are the ideal environment for the detection of cluster mergings, because the high peculiar velocities induced by the enhanced local density of the large-scale structure favour the cluster-cluster collisions, in the same way as seen in the simulations. \\\\ The Shapley Concentration supercluster represents a unique laboratory where it is possible to follow cluster mergings and to test related astrophysical consequences, as the formation of shocks, radio halos, relics and wide angle tail radiosources, and the presence of galaxies with enhanced star formation. \\\\ We present the results of an extensive multiwavelength survey of the central part of the Shapley Concentration, with the use of optical spectra, radio and X-ray data. } ",
        "introduction": "Merging clusters are the most energetic and common phenomena in the Universe, but it is not clear in which way the kinetic energy of this phenomenon is dissipated and which is the influence on the galaxy population. There are some observational features that have been associated with the cluster merging, like shocks in the hot gas, radio halos, relics and wide angle tail radiosources, and the presence of starburst galaxies, even if the precise theorical description is not yet completely assesed. \\\\ Rich superclusters are the ideal environment for the detection of cluster mergings, because the peculiar velocities induced by the enhanced local density of the large-scale structure favour the cluster-cluster and cluster-group collisions, in the same way as the caustics seen in the simulations. \\\\ The most remarkable examples of cluster merging seen at an early stage are found in the central region of the Shapley Concentration, the richest supercluster of clusters within 300 \\hmpc. This supercluster represents a unique laboratory where it is possible to follow cluster mergings and to test related astrophysical consequences. \\\\ We present the results of an extensive, multiwavelength survey of the central part of the Shapley Concentration, with the use of optical spectra, radio and X-ray data (see also the poster by Zucca et al., these proceedings). \\\\ Updates to this project can be found at the WWW page: \\\\ http://boas5.bo.astro.it/$\\sim$bardelli/shapley/shapley\\_new.html ",
        "conclusions": "In this contribution we described a multiwavelength study of two cluster complexes at the center of the Shapley Concentration: both these structures are formed by interacting clusters and have sizes of the order of $\\sim 7$ \\hmpc.  In the A3558 complex, the members of the chain appear to be fragmented in a large number of components, detected with high significance both in the optical and in the X-ray band. All these components are embedded in a hot gas filament. Moreover, Venturi et al. (2000) found a deficiency of radiogalaxies in the A3558 complex with respect to the radio-optical luminosity function of normal clusters and Bardelli et al. (1998b) found an excess of blue galaxies in the expected position of the shock. \\\\ These facts could suggest that the dynamical processes acting on this complex are in a rather advanced stage and that the merging events were already able to induce modifications in the galaxy population. Our conclusion is that the A3558 complex is a cluster-cluster collision (otherwise called ``major merging\") seen just after the first core-core encounter, where an intervening cluster impacted onto the richer object A3558. Indeed the clumpiness found eastward of A3558 could be due to the galaxies of this intervening cluster, which is now emerging from the main component. This scenario is also confirmed by numerical simulations of merging clusters (see f.i. lower panels in Fig.2 of Burns et al. 1994, which show the structure after the core-core encounter).  In the A3528 cluster complex, the effect of interactions is less spectacular. This complex appear less fragmented than the A3558 complex and the bivariate radio-optical luminosity function is not different with respect to the other clusters. \\\\ All these characteristics of the A3528 complex seem to indicate that this is a ``young\" structure (Bardelli et al. 2000b): the two interacting cluster pairs (A3528N-A3528S and A3530-A3532) resemble the simulations of Burns et al. (1994) for the pre-merger case (see upper panels of their Fig.2). This scenario is consistent with the suggestions of Reid et al. (1998), based on the analysis of extended radiosources in this region and with the fact that the merging effects on the galaxy population and on the cluster dynamics are not yet evident. \\\\ Given the overall overdensity of this region, at the end these two main components will merge together in a major merging event, forming a structure similar to the A3558 complex.  The ``age\" difference of the two complexes is mainly due to their overdensities, which lead to different collapse times (see Table 5 in Bardelli et al. 2000a). This fact confirms that the Shapley Concentration is a ``laboratory\" where it is possible to study the formation of clusters at different stages.  \\begin{figure} \\begin{center} \\epsfysize=13cm \\epsfbox{tridi.ps} \\epsfysize=8.5cm \\epsfbox{yvel.ps} \\end{center} \\par\\noindent Figure 1: Substructures in the A3528 region found in the three--dimensional sample using the DEDICA algorithm. Solid lines connect the position of each galaxy to the common limiting position $\\vec{x}_{lim}$ of the group. \\\\ Upper panel: projection on the plane of the sky. Lower panel: projection on the velocity--Y plane. \\end{figure}  \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=0.49\\hsize \\epsfbox{a3558a.ps} \\hfil \\epsfxsize=0.49\\hsize \\epsfbox{a3558b.ps} \\hfil \\leavevmode \\epsfxsize=0.49\\hsize \\epsfbox{poor.ps} \\hfil \\epsfxsize=0.49\\hsize \\epsfbox{a3556.ps} \\hfil \\end{center} \\par\\noindent Figure 2: The positions of subcluster galaxy found by DEDICA in the three-dimensional sample covering the A3558 complex are overplotted on the 2D isodensity contours of galaxy distribution. The high number of substructures reveals that this structure is far from the equilibrium state. \\\\ Upper panels: Substructures found in the A3558 cluster; for clarity, the group members have been split in two panels. \\\\ Lower panels: The same as above for SC1329-313 \\& SC1327-312 (left) and A3556 (right). \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize= 0.8 \\hsize \\epsfbox{58iso_id.ps} \\epsfxsize=0.75 \\hsize \\epsfbox{28iso_id.ps} \\end{center} \\par\\noindent Figure 3: Radio sources detected in the 22cm ATCA survey in the A3558 and A3528 complex regions overlaid on the optical isodensities. Filled circles represent optical counterparts with measured redshift, filled squares represent those without velocity information and crosses are radiosources without optical counterparts. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize= \\hsize \\epsfbox{wedge_spix.ps} \\epsfxsize= 0.9 \\hsize \\epsfbox{28wedge_id.ps} \\end{center} \\par\\noindent Figure 4: Wedge diagrams of galaxies in the A3558 and A3528 regions. Crosses are radiosources with optical counterpart. In the A3558 complex, radio galaxies with spectral index $\\alpha^{22}_{13}>1$ have been marked with a square and the extended sources have been highlighted. In the A3528 complex, extended radio galaxies are indicated by triangles. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize=0.8 \\hsize \\epsfbox{rlf_a3558.ps} \\epsfxsize=0.65  \\hsize \\epsfbox{rlf_a3528.ps} \\end{center} \\par\\noindent Figure 5: Bivariate radio-optical luminosity function for the A3558 (upper panel) and the A3528 (lower panel) complexes. Given the uncertainty in the fraction of ellipticals, the real luminosity function of the A3528 complex should lie in the shaded region. Triangles represent the normal cluster luminosity function of Ledlow \\& Owen (1996): note the lack of radio sources in the A3558 complex with respect to normal clusters. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize= \\hsize \\epsfbox{newcozza.ps} \\end{center} \\bigskip \\par\\noindent Figure 6: Radio(20cm VLA)-optical superposition at the center of the cluster A3560. This kind of information, added to X-ray maps and spectra, will constrain the dynamics of merging clusters. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize=0.75 \\hsize \\epsfbox{a3528n_new.ps} \\epsfxsize=0.75  \\hsize \\epsfbox{a3528s_new.ps} \\end{center} \\par\\noindent Figure 7: The A3528 cluster is actually a double cluster. In this figure we superimposed the radio maps of the two components (dubbed A3528N and A3528S, Schindler 1996) on the Digital Sky Survey. Note the central galaxy of A3528N, which is formed by two radiosources. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize=0.9 \\hsize \\epsfbox{scgroups_rx_h2_new.ps} \\end{center} \\bigskip \\par\\noindent Figure 8: The radio isophotes from the NVSS survey are superimposed to a part of the ROSAT PSPC image centered on A3562. At the center of the cluster A3562 (on the left) the emission of an Head Tail radio galaxy is clearly visible. Additional extended emission is present in the cluster center. At the center of the figure, an extended radio emission (possibly a mini-halo) is located at the border of the poor cluster SC 1329-313. This emission is associated to a bright cluster galaxy. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize= \\hsize \\epsfbox{ROSAT_core_new.ps} \\end{center} \\par\\noindent Figure 9: The soft X-ray emission (ROSAT mosaic from Ettori et al. 1997) in the A3558 cluster complex, with the main clusters labelled. We propose that this structure is the remnant of a cluster-cluster collision seen after the first core-core encounter. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize=0.7 \\hsize \\epsfbox{asmooth_new.ps} \\epsfxsize=0.7 \\hsize \\epsfbox{cluster_map_new.ps} \\end{center} \\par\\noindent Figure 10: Upper panel: ROSAT PSPC image of the A3562 cluster. The image has been adaptively smoothed in order to enhance the features with a significance higher than 5 $\\sigma$ with respectc to the background. Lower panel: Residuals in the emission map after subtraction of a best fit $\\beta$ model. The darker regions represent excess in emission, the lighter ones indicate a deficit. \\end{figure}  \\begin{figure} \\epsfxsize=0.7 \\hsize \\begin{center} \\rotate[r]{\\epsfbox{counts_PSPC_SAX.ps}} \\epsfxsize=0.7 \\hsize \\epsfbox{profiles_a3562.ps} \\end{center} \\par\\noindent Figure 11: Upper panel: The LECS (open squares) and PSPC (filled circles) counts of A3562 are here divided by the best-fit absorbed MEKAL model with fixed galactic absorption obtained from the MECS data (open triangles). Note that the LECS data are shifted by a factor $\\sim 0.77$ in normalization with respect to MECS and that the PSPC counts present a large disagreement below 0.5 keV. Lower panel: Projected radial temperature and abundance profiles for A3562, as determined from the BeppoSAX map. Filled circles and open triangles indicate the measurements obtained including and excluding point sources. \\end{figure}  \\begin{figure} \\begin{center} \\epsfxsize= 0.7 \\hsize \\epsfbox{SC1327_MECS_new.ps} \\epsfxsize=0.7 \\hsize \\epsfbox{SC1329_MECS_new.ps} \\end{center} \\par\\noindent Figure 12: BeppoSAX MECS maps of the two poor groups SC1327-312 and SC1329-313, which are in between A3558 and A3562. \\end{figure}"
    },
    "0011/astro-ph0011444_arXiv.txt": {
        "abstract": "Major issues and existing methods for the reduction of CMB anisotropy data are reviewed. An emphasis is put on the importance of proper modelling of the data. It is suggested that the robustness of methods could be improved by taking into account the uncertainty of the model for finding optimal solutions. ",
        "introduction": "\\label{sec:intro}  The importance of measuring anisotropies of the Cosmic Microwave Background (CMB) to constrain cosmological models is now well established. It has been demonstrated that the properties of these anisotropies depend drastically on the seeds for structure formation as well as on cosmological parameters describing the matter content, the geometry, and the evolution of our Universe \\cite{hu-sugiyama96,jungman96}.  The complexity of CMB data taking into account imperfections in the observations, however, has led the CMB community into developing sophisticated data reduction methods. In this review paper, I discuss the principal issues about CMB data analysis, centering the discussion on map-making, as other aspects related to $C_{\\ell}$ estimation and its interpretation are discussed by other authors in these proceedings. ",
        "conclusions": "In this paper, I have reviewed the main issues and methods concerning the reduction of CMB data. I have emphasised on the importance of proper modelling for finding an optimal solution to this problem. The main issues for CMB data reduction are foregrounds (astrophysical or others), slow drifts of various origins in the timelines, and instrumental imperfections as the problem of straylight pickup in the detectors. Understanding the issues and the instrumental problems is important, as the optimisation of the reduction requires optimising the data model at least as much as optimising the numerical methods for solving the problem.  Standard linear methods for handling astrophysical foregrounds and slow drifts are now well established. These methods could still be improved by taking into account explicitely the uncertainty in the measurement model (lack of knowledge of the exact number of ``sources\" for component separation, imprecision in the knowledge of foreground spectra and statistical properties, imprecision in the knowledge of detector pointings).  {\\em Acknowledgements}  Thanks to Jean-Christophe Hamilton and Alexandre Amblard for their critical reading of the original manuscript."
    },
    "0011/astro-ph0011028_arXiv.txt": {
        "abstract": "It is shown that for very rapidly rotating low mass strange stars the marginally stable orbit is located above the stellar surface. This effect is explained by the very important role of the oblateness of the rotating strange star. The comparison with some ``academic'' examples is presented. This feature is purely Newtonian in its nature and has nothing to do with relativistic marginally stable orbit. The effect is very large and cannot be treated in a perturbative way. It seems that strange stars as a very dense self-bound objects are the only possibility in Nature to represent these toy models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011502_arXiv.txt": {
        "abstract": "We discuss some of the computational challenges encountered in simulating the evolution of clusters of galaxies. Eulerian adaptive mesh refinement (AMR) techniques can successfully address these challenges but are currently being used by only a few groups. We describe our publicly available AMR code, FLASH, which uses an object-oriented framework to manage its AMR library, physics modules, and automated verification. We outline the development of the FLASH framework to include collisionless particles, permitting it to be used for cluster simulation. ",
        "introduction": "Clusters of galaxies are the largest gravitationally bound objects in the universe. They consist mainly of dark matter and diffuse, hot plasma, with galaxies themselves contributing only a few percent of the total mass. Clusters have attracted attention in recent years because they are large enough to serve as a representative sample of the universe; they provide strong contraints on cosmological models. Clusters are also interesting from an astrophysical point of view. The intracluster medium (ICM), at densities $\\sim 10^{-4}-10^{-2}$~cm$^{-3}$ and temperatures $\\sim 10^{7-8}$~K, is collisionally ionized and emits X-rays, primarily via bremsstrahlung \\cite{Sar88}. The radiative cooling time can be short enough to produce cooling flows \\cite{Fab94}. Observations of Faraday rotation show that the ICM is magnetized, with $B\\gsim 1\\ \\mu$G \\cite{Cla00}. With the diffuse, nonthermal radio \\cite{KemSar00} and X-ray \\cite{FF99} emission seen in some clusters, this suggests that clusters are sites for cosmic-ray acceleration \\cite{BBP97}. Magnetic fields may also help to suppress diffusion in the ICM \\cite{ChaCow99}. Galaxies orbiting in cluster potentials experience tidal and ram-pressure stripping and help to stir the ICM \\cite{SAP99}. Star formation and supernovae may affect the abundances of heavy elements in the ICM as well as its global energetics \\cite{DupWhi00,LPC00}. Many elements of a complete model for the ICM are known, but we still cannot answer such questions as: what is the source of entropy and metals in the ICM? what happens to the gas that cools below X-ray temperatures in cooling flows? how robust are cooling flows? how is energy partitioned among thermal and nonthermal particle populations, magnetic fields, and turbulent motions?  Cluster mergers play a key role in these phenomena. Mergers strongly affect the ICM, producing long-lived distortions in X-ray images and temperature maps. The energies they release ($\\sim 10^{64}$~erg) are easily enough to heat the ICM to $10^8$~K. These events are complex, nonlinear, multiphysical, and three-dimensional; thus numerical simulations are appropriate tools for studying them.  \\begin{figure} \\includegraphics{merger5\\paperversion} \\caption{Gas density (log contours), temperature (shading), and velocity (arrows) in a merger of clusters with mass ratio 1:3 (Ricker \\& Sarazin 2000, in prep.). Times and lengths are in Gyr and $h^{-1}$~Mpc, respectively, with Hubble constant $h=0.6$.} \\label{Fig:merger} \\end{figure}  Cluster simulations with multiphysics (at least hydrodynamics in addition to dark matter) have reached $\\gsim 10^6$ particles or zones, simulating a cluster and its environment down to kpc scales \\cite{Fre99,Lew00}. Cosmological simulations with $10^9$ particles, yielding $10^5$ clusters, have been performed \\cite{Col00}, but thus far these have only included dark matter. To understand recent observations, we must resolve scales needed for cosmological context ($\\gsim$10~Mpc) and galaxies ($\\lsim$1~kpc) -- a dynamic range of $>10^4$ -- with hydrodynamics, cooling, magnetic fields, and nonequilibrium plasma effects. Star formation and supernova feedback will remain unresolved for the forseeable future and must be included phenomenologically.  Because shocks are important in the ICM, Eulerian shock-capturing methods such as the Piecewise-Parabolic Method (PPM) \\cite{ColWoo84} are very desirable. Fig.~\\ref{Fig:merger} shows an example merger calculation performed using the COSMOS $N$-body/hydro code \\cite{RDL00}. COSMOS uses PPM on a single nonuniform grid. This 3D calculation covered a dynamic range of $\\sim 300$ on 128 processors of the San Diego Cray T3E, requiring $10,000$ node-hours. To add new physics and increase dynamic range, the cost of such calculations must be reduced significantly while retaining the shock-capturing properties of single-grid Eulerian schemes. ",
        "conclusions": ""
    },
    "0011/astro-ph0011516_arXiv.txt": {
        "abstract": "An analysis of accurate radial velocities (RVs) of the Be star \\gc from 295 Reticon spectrograms secured between October 1993 and May 2000 allowed us to prewhiten the RVs for the long-term changes and to obtain the first orbital RV curve of this star. The orbital period is 203\\D59 and the orbit has an eccentricity of 0.26. The orbital motion is detectable even in the published velocities, based on photographic spectra. This implies that \\gc is a primary component of a spectroscopic binary. The secondary has a mass of about 1~\\ms, appropriate for a white dwarf or a neutron star, but it could also be a normal late-type dwarf. The ultimate solution of the dispute whether the observed X-ray emission is associated with the secondary or with the primary will need further dedicated studies. ",
        "introduction": "The Be star \\gc~(27~Cas, HD~5394, HR~264, ADS~782A) is a member of a visual multiple system -- see also \\cite{gonc} who discovered a new companion, closer than ADS~782B. It is the very first Be star known, discovered by \\cite{secchi}. It exhibits spectral and light variations on several distinct time scales. History of its pronounced long-term spectral and light variations was summarized, e.g., by \\cite{doaz83} or \\cite{telt94}. It underwent two consecutive shell phases in 1935-36 and 1939-40, followed by a relatively short phase when it appeared as a normal B star. When the Balmer emission lines are present, they exhibit cyclic long-term $V/R$ and radial-velocity changes. In 1976, \\cite{jer} and \\cite{mason} reported that \\gc is the optical counterpart of the X-ray source MX0053+60. However, the attempts to detect the orbital motion of \\gc (on the assumption that \\gc is a binary with a compact companion), led to negative results - see \\cite{cow} and \\cite{jarad}. There is also evidence of rapid spectral changes of \\gc in the form of line-asymmetry and RV variations on a characteristic time scale of 0\\D7 -- see \\cite{john} and \\cite{jarad} -- and in the form of travelling sub-features returning every about 0\\D8 to the line centre  -- see \\cite{ninkov, yang, hora} and \\cite{smith95}. Later, \\cite{smith1} found simultaneous but anti-correlated UV and X-ray light changes with a period of 1\\D123 which they identified with the rotational period of \\gce. Note, however, that \\cite{tow3}, searching for such a period only in a close neighbourhood of the 1\\D123 period, found a significantly different period of 1\\D15655\\p0\\D00012 from the Hipparcos $H_p$ band photometry. The complex nature of rapid line-profile changes of the UV Si~IV lines was demonstrated by \\cite{smith2}.  There are currently two competing interpretations of the nature of the observed X-ray emission: one is the accretion of the wind from \\gc onto a putative white-dwarf companion and the other one is that it originates from some physical processes in the outer atmosphere of \\gc itself. Arguments for and against these two hypotheses are best summarized in recent studies by \\cite{kubo} and \\cite{rob2000}.  Realizing that there had been no attempt to search for the orbital motion of \\gc in electronic spectra, we have been collecting Reticon spectrograms of \\gc since 1993. Here, we report the first results.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{ff1.eps}} \\caption[ ]{The \\ha and He~I~6678 line profiles from an Ond\\v{r}ejov Reticon spectrum of \\gc obtained on HJD~2449310.1969. Note the different intensity scale at the two panels}\\label{sp} \\end{figure} ",
        "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{ff5.eps}} \\caption[ ]{A plot of Hipparcos photometry, prewhitened for a linear trend, vs. phase of the best-fit short period of 1\\D48700. Phases are calculated from the epoch of minimum light, HJD~2448350.838. To illustrate the noise of Hipparcos photometry, the original $H_p$ photometry of a constant 5-mag. B star 4~Per, suitably shifted in magnitude, is also shown}\\label{hipv} \\end{figure}  The latest arguments by \\cite{rob2000} against the origin of the X-ray emission from an accretion onto a white-dwarf companion appear quite convincing. However, the understanding of the physical processes involved is still far from complete. We mention two possibly relevant items: (i) We found that \\gc is a binary with a 203\\D59 period and that its orbit is eccentric ($e=0.26$), similarly as for other Be+X binaries. Other known Be binaries usually have circular orbits. The varying distance between the components should be seen to affect the putative accretion onto the secondary. A search for possible modulation of the X-ray flux with the phase of the orbital period is, therefore, highly desirable. (ii) Considering the importance of the 1\\D12 period for the alternative hypothesis, we re-analyzed good (flag = 0) Hipparcos $H_p$ photometry of \\gce, prewhitened for a linear trend, over a wider range of possible short periods than \\cite{tow3} did: between 0\\D6 and 5\\D0. The best-fit period is 1\\D48700\\p0\\D00013, with minimum light at HJD~2448350.838\\p0.027 -- see Fig.~\\ref{hipv} which shows the superiority of this period to the period found by \\cite{tow3}. The new period is very significantly different from 1\\D12. If real, it could be either a corotation period at some level or a low-order mode of pulsation.  It is, therefore, desirable to carry out further tests of both competing hypotheses."
    },
    "0011/astro-ph0011499_arXiv.txt": {
        "abstract": "\\index{abstract} We report results from the first 1000 square arc-minutes of the Las Campanas IR survey. We have imaged 1 square degree of high latitude sky in six distinct fields to a 5$\\sigma$ H-band depth of 20.5 (Vega). Optical imaging in the V,R,I,and z' bands allow us to select color subsets and photometric-redshift-defined shells. We show that the angular clustering of faint red galaxies ($18 < H < 20.5, I - H > 3$) is an order of magnitude stronger than that of the complete H-selected field sample. We employ three approaches to estimate $n(z)$ in order to invert $w(\\theta)$ to derive $r_0$. We find that our $n(z)$ is well described by a Gaussian with $ <z> = 1.2, \\sigma(z) = 0.15$. From this we derive a value for $r_0$ of $7^{+2}_{-1}$~ co-moving $h^{-1}$Mpc at $<z> = 1.2$. This is a factor of $\\sim 2$ larger than the clustering length for Lyman break galaxies and is similar to the expectation for early type galaxies at this epoch. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011450_arXiv.txt": {
        "abstract": "Quantification of the Galaxy's star formation history involves both the duration and the rate of formation, with these parameters being known with different precision for different populations. The early rate of star formation is knowable from modelling chemical element data, the recent rate directly from isochrone analyses of colour-magnitude data. The field halo and globular clusters are almost exclusively old, and formed in at most a few Gyr. The outer bulge probably formed in a short period long ago -- extant data is inconsistent, while the inner bulge/disk is forming today, and has continued to form over time. Only very limited data is available on the inner disk. The outer disk near the Sun seems as old as the halo. The earliest extended disk, which forms the thick disk today, seems to have been in place very early, an observation which is not simply consistent with some galaxy formation models. ",
        "introduction": "Which is more important: do galaxies form, or are they assembled? More explicitely, does star formation occur primarily in the eventual gravitational potential, albeit in localised regions within that potential, or does it occur in regions with much smaller potential wells, which are later assembled in the current whole?  The star formation history of a galaxy, explicitely here our Milky Way Galaxy, where the most detailed information is attainable, is the convolution of two functions. One function describes the rate of formation of the stars which are today in the Galaxy. The second describes the assembly of those stars into the present Galactic potential well. There is direct evidence that this assembly continues today, with both stars and gas being assembled into, or at least rearranged in, the Galactic potential (eg, the Sagittarius dwarf spheroidal, Ibata, Gilmore \\& Irwin 1994, 1995; the Magellanic Stream, Putman etal 1998). HST imaging suggests that the rate of accretion/merging was significantly higher in the past, so that accretion has always been significant. However, HST and other data also show that the star formation rate was higher in the past.  A topical question is then the relative importance of these two processes, the normalisation of the merger and star formation functions, and the order in which they occur: does star formation mostly happen during mergers, or after two potential wells have merged?  Some fairly direct constraints on the relative importance of star formation in a galaxy-scale potential can be deduced from the galaxian luminosity metallicity relation, shown in figure~1 for the Local Group. For Local Group galaxies direct metallicities can be determined, obviating possibly model-dependent interpretation of line indices for unresolved populations. Nonetheless, this relation is consistent with the more general mass-metallicity relation for galaxies, suggesting this relation is universally valid, and applicable over the full mass range of galaxies.  \\begin{figure}[ht] \\plotone{gilmoreg_fig1.ps} \\caption{The luminosity-metallicity relation for Local group galaxies. The rather good correlation indicates that the average star forms in a potential well which is related to the potential well today, even though the mean stellar age in these galaxies covers a wide range. Equivalently, one cannot build a large galaxy by assembling smaller stellar ones.}  \\end{figure}    The mass-metallicity relation shown in Fig~1 indicates that when the typical star in a galaxy is formed, at whatever redshift, it knows the depth of the gravitational potential in which it will be orbiting at redshift zero. This knowledge indicates that one cannot form a large galaxy by assembling smaller stellar systems: any such resultant object would be of high luminosity and stellar mass, but of low metallicity. Such objects are not found.  The natural solution is to form most stars, and most metals, either during mergers, or later after assembly of gas. Star formation during mergers of comparable mass gassy systems, while a very visible process which is clearly important still today, and a natural way to build bulges, is a natural way to build a disk galaxy only if star formation is very inefficient {\\sl during} the merger. It is also far from a natural way to build a very low-mass galaxy.  One might be tempted to conclude that low mass galaxies and galaxy disks formed their stars not only {\\sl in situ}, but at a lower rate than was relevant to bulge formation.  We now consider local constraints on this speculation.  \\subsection{ Global considerations}  The rate at which stars formed, on average in tolerably luminous systems, has been recently quantified in the Madau plot. In linear time, rather than redshift space, this shows a roughly constant rate of star formation from very early times ($\\la 1$Gyr afer time zero) continuing for 4$-$5Gyr, then an apparently rapid decline (no doubt exaggerated by the non-linear time-redshift relation) by a factor of about 4$-$5 to a new roughly constant rate, continung for a further 5$-$6 Gyr until today. Thus, some two-thirds of stars were formed before redshift about unity. Does the Local Group follow this trend? If it does, then both the whole of their bulges, and the inner disks of the Galaxy and M31 must have been in place, and stellar, before redshift unity.  We discuss this further below, but note here that the extant, albeit indirect evidence, suggests consistency with the Madau plot.  The most directly consistent interpretation is that the bulge and early disk (now the thick disk) were formed in the first $\\sim 2-3$Gyr, at redshifts significantly greater than unity. ",
        "conclusions": ""
    },
    "0011/astro-ph0011100_arXiv.txt": {
        "abstract": "s{ We report the results of spatially resolved X-ray spectroscopy of 8 different {\\sl ASCA} pointings distributed symmetrically around the center of the Perseus cluster. The outer region of the intracluster gas is roughly isothermal, with temperature $\\sim$ 6--7 keV, and metal abundance $\\sim$ 0.3 Solar. Spectral analysis of the central pointing is consistent with the presence of a cooling flow and a central metal abundance gradient. A significant velocity gradient is found along an axis highly discrepant with the major axis of the X-ray elongation. The radial velocity difference is found to be greater than 1000 km s$^{-1}Mpc^{-1}$ at the 90\\% confidence level. Simultaneous fittings of GIS 2 \\& 3 indicate that two symmetrically opposed regions have different radial velocities at the 95\\% confidence level and the F-test rules out constant velocities for these regions at the 99\\% level. Intrinsic short and long term variations of gain are unlikely (P $<$ 0.03) to explain the velocity discrepancies. We also report the preliminary results of a similar analysis carried out for the Centaurus cluster, where long-exposure SIS data is available. We also find a significant velocity gradient near the central regions (3$^\\prime$-8$^\\prime$ of Centaurus. If attributed to bulk rotation the correspondent circular velocity is $\\sim$1500$\\pm$150~km s$^{-1}$ (at 90\\% confidence). The line of maximum velocity gradient in Centaurus is near-perpendicular to the infalling galaxy group associated with NGC 4709.} ",
        "introduction": "X-ray determination of the physical state of the intracluster gas provides a unique tool to probe the origin and evolution of clusters of galaxies.  {\\sl ASCA} observations have shown significant spatial temperature variations in many clusters, indicating that clusters are currently evolving systems. This is consistent with the predictions of hierarchical cluster formation within CDM models, and had been suggested in pre-{\\sl ASCA} times \\cite{fit,u92}.  Although early models of galaxy clusters treated them as spherically symmetric virialized systems, recent X-ray and optical studies show that often clusters show substructures. Furthermore, if cold dark matter models (CDM) are correct, primordial small-scale density fluctuations are not erased and clusters are formed by infall/merging of smaller scale objects (bottom-up hierarchical scenario). The merger (infall) of sub-clumps creates shocks (associated with temperature substructure), bulk gas flows (associated with velocity substructure) and asymmetric distributions of velocity distribution of galaxies \\cite{bird}. In order to understand the physical properties of clusters, their origin and evolution one has to take into account the degree and the physical scale of substructuring.  The determination of complex temperature substructure in clusters is often interpreted to be related to shocks due to cluster merger at different stages. The link between temperature substructure and the merger stage is often done by comparison with hydrodynamical simulations. There is currently an enormous variety (different initial conditions, hydro-codes, impact parameters, matter components, etc.) of cluster formation/merger simulations in the literature \\cite{ev,kw,rbl93,rbl97,sch,ptc,nfw,rlb97,emn,rick,tak98,tak99,tak00}, which provide simulated temperature maps that can be used to compare with the observations. Temperature maps provide important clues about the merging state of a cluster and the degree of substructuring, yet a comparison between temperature maps and numerical simulations is likely to show multiple configurations that may explain equally well what is observed, depending on the line of sight the observer chooses and also on the resolution of temperature maps, especially when one considers off-center mergers \\cite{tak00,rick,rof}. This is because we are always looking at 2-D projections of the cluster, so additional information is needed descibe the cluster evolution stage accurately (e.g. gas velocity distribution).  A more straightforward way to determine the effects of a merger is to directly measure intracluster gas velocities. Several simulations of cluster mergers \\cite{rick,rbl93,tak99,rof} indicate residual gas velocities of a few thousand km~s$^{-1}$. To measure gas velocities one requires accurate determinations of spectral line centroids. The precision with which a line centroid can be measured, in velocity space, is $\\sim 127~\\Gamma_{eV}(E_{keV} N^{\\frac{1}{2}})^{-1}$km~s$^{-1}$, where N is the number of photons in the line and $\\Gamma_{eV}$ is the FWHM of the line, or if the line is narrower than the instrumental width, is the FWHM of the instrument, and E$_{keV}$ is the line energy. The energy resolution of the spectrometers on-board {\\sl ASCA} vary from 2\\% (SIS) to 8\\% (GIS) at 5.9 keV. For a FWHM of 9000 km~$s^{-1}$ at 6.7 keV,  which is typical of early (first 3 years) SIS data at the FeK$\\alpha$ line, to obtain a line centroid to a precision of 500 km/s, we need $\\sim$ 60 line photons.  This same accuracy can be obtained with the GIS with $\\sim$ 350 line photons. Since these estimates become more uncertain if the rms of the continuum starts to compete with the line, and, to compensate for the uncertainties introduced by {\\sl ASCA} PSF, one would like to analyze clusters that have high metal abundance and several different long exposures observations from off-center regions surrounding the X-ray center. Perseus (Abell 426) and Centaurus (Abell 3526) clusters match this criteria and, therefore, are well suited for the study of gas velocity distribution. ",
        "conclusions": "The spectral analysis of {\\sl ASCA} data carried out in this work indicates the presence of a significant velocity gradient in the ICM of the Perseus \\& Centaurus clusters. In Perseus two regions symmetrical with respect to the cluster center show discrepant redshifts not just with respect to each other but also with respect to the central region. This velocity difference is unlikely to be attributed purely to gain fluctuations and suggests the existence of large-scale bulk motions of the intracluster gas. The two symmetrically opposed discrepant regions have velocity differences of $\\sim$~-3000 km~s$^{-1}$ (or $\\le$ -~600 km~s$^{-1}$ at the 90\\% confidence level) for P1 and $\\sim$ +5000 km~s$^{-1}$ (or $\\ge$ +~60 km~s$^{-1}$ at the 90\\% confidence level) for P5 with respect to P0. There are no observed temperature or abundance differences between the two pointings with discrepant redshifts (P1 \\& P5). The velocities measured are consistent with large scale gas rotation with a correspondent circular velocity of $\\sim$ 4100 $^{+2200}_{-3100}$ km s$^{-1}$ (90\\% confidence) \\cite{db1}. This implies a large angular momentum for the ICM and that a significant fraction of the gas energy may be kinetic.  In Centaurus, where analysis of SIS data allow more precise velocity measurements, we also find evidence for significant velocity differences consistent with bulk rotation with circular velocity of $\\sim$ 1500$\\pm$150~km s$^{-1}$ (at the 90\\% confidence) at 5$^\\prime$ from the cluster center. There is also evidence for an azimuthal temperature gradient in Centaurus, where the temperature grows from 2.73$\\pm$0.06 keV in the northwestern direction to 3.6$\\pm$0.11 in the Southeastern region, towards the apparently merging galaxy group centered in NGC 4709.  The best candidate for generating this large angular momentum in both clusters is off-center mergers. In off-center mergers, up to $\\sim$ 30\\% of the total merger energy may be kinetic (can be transferred to rotation) \\cite{ptc}. Off-center merger simulations \\cite{rick,tak99,tak00,rof} often produce residual intracluster gas rotation with velocities of a few thousand km~s$^{-1}$. Additional evidence for merging comes from the observed offset between the optical center and the X-ray center (Perseus) \\cite{bran,sny,u92}; the radial change in X-ray isophotal orientation (isophotal twist) \\cite{mohr}; the asymmetric galaxy morphological distribution, with preferential eastward direction in the distribution of E+S0 types (Perseus) \\cite{bru}; the bimodal velocity distribution coincident with temperature substructure (Centaurus) \\cite{chu,lucey}. However, simulations also indicate other observable consequences of mergers that can, in principle, be cross-analyzed with the velocity maps to test the robustness of the merger scenario. One of the features predicted by off-center cluster-cluster mergers is a strong negative radial temperature gradient (core heating) ($\\ge$ 2 keV/Mpc) for most of the merger life-time, even when the angle of view is not favored, e.g. along the collision axis \\cite{tak00,rick}. In our case we do not observe a negative temperature gradient at all, actually we observe a positive temperature gradient in both clusters due to the cooling flow. The mere fact that the cooling flow is present makes the off-center merger explanation more uncertain, since it has been suggested \\cite{esf,rbl93} that a merger would disrupt any pre-existing cooling flows. However, recent simulations of head-on cluster mergers indicate that cooling flows can survive mergers depending on the produced ram-pressure of the gas in the infalling cluster \\cite{gomez}. Even if cooling flows are disrupted by a merger they can be reestablished quickly if the cooling time of the primary pre-merger component is small \\cite{gomez}. The fact that the major axis of the X-ray elongation is relatively close to the apparent rotation axis is another difficulty posed to the off-merger explanation. In most cases simulations show that the isodensity contours are elongated perpendicularly to rotation axis, except in some short-lived merger stages viewed from specific directions \\cite{tak00}.  The velocity and temperature distributions of Centaurus allow for the possibility that we are seeing a pre-merger stage with the (NGC 4709) sub-group infalling with a significant velocity component on the line of sight and on the E-W directions. In the case of Perseus, however, a pre-merger scenario does not fit the observations easily. This is, in part, because we do not observe X-ray surface brightness enhancement at the direction of the redshift-discrepant regions and also because there are no clear temperature substructures associated with the regions of discrepant velocities \\footnote{We do observe a mild enhancement in surface brightness towards the East region of Perseus, coinciding with our P4 pointing. This enhancement was noticed previously \\cite{schw,ettori} with better significance, and has been interpreted as evidence for merging. Although there is no clear optical evidence of a sub-group towards the direction of P5 that corroborates this scenario, the region of surface brightness enhancement (P4) seems to be associated with a higher fraction of early-type galaxies \\cite{bru}.}. Given that in some off-center merger simulations high rotational gas bulk velocities can be maintained for $\\ge$ 3 crossing times \\cite{rick} and that cooling flow survival to mergers is more likely to happen in off-center mergers \\cite{gomez} the results are consistent with a large off-center merger event in Perseus $\\ge$ 4 Gyr ago (assuming a cluster mass \\cite{ettori} of 5$\\times$10$^{14}$ M$_{\\odot}$ at a radius of 1 Mpc and that the rotating gas at $\\sim$ 7 KeV is gravitationally bound).  Velocity measurements of intracluster gas with {\\sl Chandra} and, especially, {\\sl XMM-Newton} satellites will be able determine ICM velocities in the central regions more precisely, thus providing information on the gas velocity curve, which will strongly constraint cluster-cluster merger models, or suggest alternatives for generating the large angular momentum observed."
    },
    "0011/astro-ph0011336_arXiv.txt": {
        "abstract": "The Accretion-Ejection Instability (AEI) is explored numerically using a global 2d model of the inner region of a magnetised accretion disk.  The disk is initially currentless but threaded by an external vertical magnetic field created by external currents, and frozen in the flow.  In agreement with the theory a spiral instability, similar in many ways to those observed in self-gravitating disks, but driven by magnetic stresses, develops when the magnetic field is close to equipartition with the disk thermal pressure.  The present non-linear simulations give good evidence that such an instability can occur in the inner region of accretion disks. ",
        "introduction": "The accretion-ejection instability (AEI) (Tagger \\& Pellat \\cite{TaggerP99}; see also Varni\\`ere et al., Rodriguez et al., and Tagger et al., this workshop) can occur close to the inner edge of a magnetised accretion disk when the plasma beta (the ratio of gas to magnetic pressures) is around unity. The instability appears, if the magnetic field decreases outwards sufficiently fast, as a spiral wave of low azimuthal wavenumber, made unstable through the interaction with a Rossby vortex (associated with the gradient of vorticity) which it generates at its corotation radius.  We have undertaken numerical simulations, with a setup optimised according to our previous knowledge of the AEI. We consider an infinitely thin disk threaded by a moderate vertical magnetic field: this is justified by the properties of the AEI, which are essentially constant vertically across the disk (unlike the magneto-rotational instability (Balbus \\& Hawley, 1991) for which vertical modes are required).  The disk is embedded in vacuum and, in order to separate different physical effects, we consider here only configurations where initially the equilibrium magnetic field is due to external currents. Starting with small scale initial random fluctuations in the fluid velocity, the simulations show large-scale spiral perturbations evolving in the disk, and corresponding to the expected properties of the AEI.  The following sections present the essentials of the model, basic results and conclusions. More detailed discussions of these are covered elsewhere (Caunt \\& Tagger, 2000). ",
        "conclusions": "Using this model we have shown the essential features predicted by linear theory of the AEI. We have seen that the instability is clearly strongest for field strengths of the order of equipartition, and leads to the development of low azimuthal wavenumber spiral waves and a Rossby vortex.  Future work will link in the third type of wave (namely the vertical emission of Alfv\\'en waves) from which the formation of winds and jets is believed to be possible."
    },
    "0011/astro-ph0011046_arXiv.txt": {
        "abstract": " ",
        "introduction": "Blazars, including BL Lacertae objects and Optical Violently Variable (OVV) quasars, are a class of active galactic nuclei of which the spectrum is dominated by nonthermal radiation from relativistic electrons in jets pointing at us (e.g. Blandford, Rees 1978; Blandford, K\\\"{o}nigl 1979). Their multi-frequency spectra exhibit two pronounced components. According to the unified model of blazars, the low energy component arises via synchrotron radiation (SR), and the high energy one via inverse-Compton (IC) scattering by the same population of electrons that produce the SR component. Seed soft photons for the IC scattering can be either the SR photons themselves (so-called synchrotron-self-Compton process; Maraschi, Ghisellini, Celloti 1992; Bloom, Marsher 1996), ultraviolet photons from the accretion disk (e.g. Dermer, Schlickheiser, Mastichiadis  1992) or from emission line cloud (Sikora, Begelman, Rees 1994; Blandford, Levinson 1995), or infrared photons ambient to the host galaxy (Sikora et al. 1994).  There are two extreme subclasses of blazars (Sambruna, Maraschi, Urry 1996; Giommi, Ansali, Micol 1995; Fossati et al. 1998). One of them, high-energy peaked BL Lac objects (HBLs; Padovani, Giommi 1996), have the SR peak in the ultraviolet to soft X-ray band. Their X-ray emission is dominated by the SR component as evidence by steep spectral slopes with photon index $\\Gamma=2.0\\sim3.0$ (e.g. Kubo et al. 1998). The other, OVV quasars, show the SR peak in the millimetor to optical range, and their X-ray emission is dominated by the IC component with relatively flat spectral slopes with $\\Gamma\\sim1.5$ (Kubo et al. 1998).  Yet, another subclass of blazars called low-energy peaked BL Lac objects (LBLs; Padovani, Giommi 1996) exhibit spectral properties intermediate between those of HBLs and OVV quasars (Sambruna, Maraschi, Urry 1996; Giommi, Ansali, Micol 1995; Fossati et al. 1998). In their spectra, the SR and IC components are thought to compete in the ultraviolet to X-ray bands, resulting in X-ray slopes of $\\Gamma=1.5\\sim2.0$ (Kubo et al. 1998). In order to disentangle interplay between the SR and IC components, and to examine the two component interpretation of the blazars emission, extensive X-ray observations of LBLs have been carried out (e.g. Urry et al. 1996; PKS 0735+178, Madejski, Schwartz 1988; PKS 0521-365, Garilli, Maccagni 1990; BL Lacertae itself, Kawai et al. 1991; AO 0235+164, Madejski et al. 1996). However, there are only a few LBLs of which the X-ray spectra have been decomposed unambiguously  into the two components.  OJ 287, at a redshift of $z=0.306$ (Stickel et al. 1989), is one of the most extensively observed BL Lac objects. It exhibits large intensity and polarization variability in various energy bands, together with superluminal motion (Gabuzda et al. 1989). Its multi-frequency continuum is smoothly distributed from radio to ultraviolet bands, with a SR peak at  $\\sim 5 \\times 10^{13}$ Hz and radio to X-ray spectral index of $\\alpha_{\\rm rX}\\sim0.84$ (Sambruna, Maraschi, Urry 1996); these make OJ 287 a typical LBL (Padovani, Giommi 1995). In X-rays, relatively steep spectra were obtained with Einstein, EXOSAT and ROSAT in 1979 -- 1980 (Madejski, Schwartz 1988), 1983 -- 1984 (Sambruna et al. 1994) and 1991 (Comastri, Nolendi, Ghisellini 1995; Urry et al. 1996), respectively. However, the Ginga observations performed in 1989 -- 1990 (Ohashi 1989; Tashiro 1994) failed to detect X-rays from this object. The X-ray intensity varied by a factor of 3 on time scales of months (Madejski, Schwartz 1988), and 30\\% within 3 days (Pollock et al. 1985). The X-ray spectra obtained with the Einstein IPC, smoothly connecting to the lower energy continua (Madejski, Schwartz 1988), is considered to arise from the SR process. Giommi et al. (1995), however, detected a hint of X-ray excess above extrapolation from the  optical-to-ultraviolet continua. This suggests the presence of an additional IC component in the X-ray emission from this  LBL.  In 1994, an optical outburst of OJ 287 motivated a series of worldwide multi-band observations including an ASCA observation. However, the measured X-ray flux was the lowest among the X-ray results reported so far (Idesawa et al. 1997; paper I). Moreover, the obtained X-ray spectra, exhibiting the hardest record, did not smoothly connect to the low energy continuum. As argued by Idesawa et al. (paper I), this reinforce the presence of the underlying IC component in the ASCA X-ray spectra, although they could not unambiguously resolve the suggested IC component from the SR component. In order to better resolve the two components, it is important to measure long-term spectral and intensity changes of this object, where the two emission components are expected to vary independently to a first approximation. Accordingly, two separate ASCA observations were conducted in 1997. ",
        "conclusions": "The BL Lacertae object  OJ 287 was observed twice in 1997 with ASCA, while the source was getting optically fainter after the optical outburst in 1994. The X-ray spectra obtained with the two sets of focal-plane instruments on board ASCA have been consistently  described with a single  power-law model modified with the Galactic absorption in both observations. The obtained source flux densities at 1 keV are lower and the spectra are harder % in comparison with the previous X-ray results. No significant evidence of short-time variability was observed on time scales of hours.  Figure 3 shows the multi-frequency spectra of OJ 287 obtained in 1994 and 1997, including the present X-ray results. Radio and optical data in 1997 are presented by courtesy of Dr. Tapio Pursimo and the X-ray data in 1994 are taken from paper I. Compared with smooth extrapolations from the lower frequency continua, the X-ray spectra obtained with ASCA exhibit higher fluxes and flatter slopes in both observation results. Therefore the X-rays cannot be explained as an extension from the lower frequency SR component. Furthermore, the derived photon indices of $\\Gamma\\sim1.5$, which are consistent with those of other well studied LBLs (Kubo et al. 1998), are much flatter than those of the SR component in well studied HBLs, typically $\\Gamma=2.0\\sim3.0$ (Sambruna et al. 1996; Kubo et al. 1998). On the other hand, these photon indices are closer to those of IC X-rays observed from OVV quasars (e.g. Makino et al. 1989, 1991; Tashiro 1994; Kubo et al. 1998). These results suggest that the X-rays from OJ 287 observed with ASCA are dominated by the IC component, which is produced by electrons responsible for the radio to optical SR component.  Figure 4 shows the behavior of OJ 287 on the plane of X-ray photon index versus X-ray flux density at 1 keV ($F_{1{\\rm keV}}$), based on  available measurements taken from paper I. Comparing the Einstein, EXOSAT, and ASCA measurements covering similar energy bands, the plot shows a constant photon index $\\Gamma\\sim1.5$ in $F_{1{\\rm keV}}<0.5~\\mu$Jy, while the photon index increases with flux for $F_{1{\\rm keV}}>1~\\mu$Jy. Similar spectral behavior can be seen in some well-studied LBLs, such as S5 0716+714 (Giommi et al. 1999), ON 231 (Tagliaferri et al. 2000), and BL Lacertae (Tanihata et al. 2000), and is opposite to those of SR X-rays observed from HBLs which exhibit a harder X-ray spectrum for increasing intensity. This behavior of OJ 287 is interpreted that a decrease in the softer  SR component unveils the harder IC components that is less variable. The systematic difference between the ROSAT data points and the ASCA-Einstein-EXOSAT ones can be explained that ROSAT, with its lower energy band, is more sensitive to the SR component than the other experiments. The convergence of figure 4 at photon index $\\Gamma\\sim1.5$ and the stable light curve on short timescale  strongly suggest that the X-rays observed in 1997 consist almost solely of the IC component.  Since the X-ray spectrum observed in the 1994 optical outburst is somewhat steeper ($\\Gamma\\sim1.7$; paper I) than the limiting value of $\\Gamma\\sim1.5$, it is thought to be a mixture of the SR and the IC components. In fact, a re-analysis of the 1994 data reveals that the spectrum above 2 keV is well described with a single power law model of $\\Gamma=1.54\\pm0.04$ modified with the Galactic absorption ($\\chi^2/{\\rm d.o.f.}=209.6/232$). This photon index is consistent with that obtained on 1997 November 18, when the emission is almost purely of IC origin. However, using this parameter, the spectrum in the full energy range of ASCA ($0.8\\sim10$ keV and $0.6\\sim7.0$ keV for the GIS and SIS, respectively) shows a soft excess below 2 keV ($\\chi^2/{\\rm d.o.f.}=540.0/369$). Thus, the 1994 spectrum is examined against a double power-law model modified with the Galactic absorption. One power-law represents the IC component, with photon index fixed at the best fit value of 1.51 in November 1997, and normalization also fixed at the value determined by the single power law fit above 2 keV. The other power-law represents the SR component, whose spectral parameters are left free. This model has given an acceptable fit ($\\chi^2/{\\rm d.o.f.}=368.1/368$) to the 1994 ASCA spectra, yielding the parameters given in table 2. In addition, the estimated photon index ($\\Gamma \\sim 2.6$) for the SR component is reasonable in comparison with the observed SR components in HBLs (Sambruna et al. 1996; Kubo et al. 1998). Figure 5 shows the SR and IC components obtained in this way, together with the multi-frequency spectrum observed in 1994. When  extrapolated to lower frequencies, the SR component smoothly connects to the observed lower frequency SR continua.  In conclusion, as suggested by Idesawa et al. (paper I), the X-ray spectrum of OJ 287  is explained as a composite between the SR and the IC components. The present X-ray emission observed in 1997 is naturally attributed almost entirely to the IC emission, while the 1994 spectrum is partly contributed by the SR component. These results give a support to the unified scheme of blazars, in which their multi frequency spectra are decomposed into SR and IC components, which compete in the X-ray band for LBLs.\\par \\vspace{1pc}\\par We thank Dr. Tapio Pursimo for providing us with the Radio/optical/ultraviolet intensity data in advance of publication. We also thank all the members of the ASCA team for mission management, including spacecraft operation, data acquisition and instrumental calibrations."
    },
    "0011/astro-ph0011270_arXiv.txt": {
        "abstract": "We present low- and medium resolution spectra of the recurrent nova CI~Aquilae taken at 14 epochs in May and June, 2000. The overall appearance is similar to other U~Sco-type recurrent novae (U~Sco, V394 CrA). Medium resolution ($\\lambda/\\Delta\\lambda\\approx7000-10000$) hydrogen and iron profiles suggest an early expansion velocity of 2000--2500 km~s$^{-1}$. The H$\\alpha$ evolution is followed from $\\Delta$t = $-$0.6 d to +53 d, starting from a nearly Gaussian shape to a double peaked profile through strong P-Cyg profiles. The interstellar component of the sodium D line and two diffuse interstellar bands put constraints on the interstellar reddening which is estimated to be E(B$-$V)=$0\\fm85\\pm0\\fm3$. The available visual and CCD-V observations are used to determine t$_{\\rm 0}$, t$_{\\rm 2}$ and t$_{\\rm 3}$. The resulting parameters are: t$_{\\rm 0}$=2451669.5$\\pm$0.1, t$_{\\rm 2}=30\\pm1$ d, t$_{\\rm 3}=36\\pm1$ d. The recent lightcurve is found to be generally similar to that observed in 1917 with departures as large as 1--2 mag in certain phases. This behaviour is also typical for the U~Sco subclass. ",
        "introduction": "The first documented outburst of CI~Aql was discovered in 1917 at a photographic maximum of 11\\fm0 (Reinmuth 1925). The reported $\\sim4\\fm6$ eruption and the lack of subsequent outbursts suggested CI~Aql to be a classical nova or dwarf nova with long recurrence time. The dwarf nova classification was disproved by spectroscopic observations which showed no Balmer emission lines (Szkody \\& Howell 1992). Very recently, Wilson (2000) examined the Harvard College Observatory patrol plates taken in 1917 and found the 1917 outburst much brighter than Reinmuth's estimate. Tycho B magnitudes for nearby comparison stars gave a maximum brightness of 8\\fm6.  Time resolved CCD photometric observations by Mennickent \\& Honeycutt (1995) lead to the discovery of eclipsing binary nature of CI~Aql with an orbital period of 0.618355 days. Their subsequent optical and far red spectroscopy showed only absorption lines of hydrogen, sodium and calcium. The Ca II infrared triplet being a considerably good luminosity indicator implied an evolved giant secondary component and the authors suggested CI~Aql to belong to the class of recurrent novae. However, the lack of short-term erratic light variation (flickering) lead to a conclusion that the components of CI~Aql were not in strong interaction between 1991-1995. In contrast to above, Greiner et al. (1996) presented optical spectra in the 3800--7200 \\AA\\ range showing the high-excitation lines HeII $\\lambda4686$ and NIII/CIII $\\lambda4650$ in emission. Also, they found the [O I] emission at 6300 \\AA. Based on a simple consideration about the mass dependence of the semi-major axis for a {\\it reasonable} range of masses of the two components, Greiner et al. excluded the possiblity of a giant component in the system, but suggesting subgiant or main-sequence components. However, their results were inconclusive on the real nature of CI~Aql. The present outburst cleared the situation, placing CI~Aql among the recurrent novae.  Recurrent novae (RNe) form a rare class of cataclysmic variables. Their repeated outbursts are thought to be caused by thermonuclear runaways on the surface of white dwarf primaries which accrete material from the cool dwarf or giant secondary stars. Three subclasses (T~Pyx, U~Sco, and T~CrB) can be distinguished according to the physical nature of the secondary and the corresponding orbital period (Warner 1995). The 0.618 d orbital period of CI~Aql is typical for the U~Sco subclass (V394~CrA, LMC-RN, U~Sco) consisting of RNe with He-dominated quiescent disks, $P_{\\rm orb}\\sim 1$ d and visible secondary spectra in quiescence. The declining branches of their lightcurves vary from outburst to outburst indicating intrinsic variations of the interaction between the components.  The 2000 outburst of CI~Aql was discovered by K. Takamizawa on unfiltered T-Max 400 exposures taken on April 28.669 and 28.673 UT. An independent photographic discovery was reported by M. Yamamoto. The first low-resolution spectra (April 29.6 UT) showed prominent H$\\alpha$ emission with no P-Cyg profile indicating a fast nova at a slightly evolved stage (Takamizawa et al. 2000). Precise astrometry by Yamaoka et al. (2000) resulted in the identification of the star with CI~Aql known as an eclipsing binary (Mennickent \\& Honeycutt 1995) and Nova~Aql 1917 (Reinmuth 1925, Duerbeck 1987). Their coincidence was provided by the large scale of the discovery plates taken in 1917. Objective prism spectra by W. Liller showed strong H$\\alpha$ emission, weak H$\\beta$ line, [O I] lines (6300 and 6360 \\AA) and He II at 5870 \\AA (Yamaoka et al. 2000). Early multicolour photometric data obtained around maximum (Yamaoka et al. 2000, Wilson et al. 2000, Hanzl 2000) gave (B$-$V)$_{\\rm max}\\approx0\\fm7-0\\fm8$ suggesting relatively high reddening, also pointed out by Liller based on the observed weak H$\\beta$ emission compared to the strong H$\\alpha$ line. Mazuk et al. (2000) reported 0.8--2.5 micron spectrophotometry obtained 74 days after the peak brightness, which showed evidence for an increased excitation of the emission-line gas, as the neutral C, N and O lines disappeared and the prominent He II lines replaced the He I 10830 \\AA\\ line.  The main aim of this paper is to present optical spectra taken around and after the maximum, between $\\Delta$t = $-$0.6 d to $\\Delta$t = +53 d. The low- and medium resolution spectra were used to determine the main outburst properties, the expansion velocity and the interstellar reddening. We also use all publicly available visual photometric data collected by the VSNET group ({\\tt http://www.kusastro.kyoto-u.ac.jp/vsnet}) to estime the rates of decline and to check the photometric phases of the obtained spectra.  The paper is organised as follows. Observations and data reductions are briefly described in Sect.\\ 2, the detailed description of the obtained spectra is given in Sect.\\ 3, while the reddening determination is discussed in Sect.\\ 4. We present the visual lightcurve in Sect.\\ 5 and compare it with the 1917 outburst. Concluding remarks are given in Sect.\\ 6. ",
        "conclusions": "The presented spectroscopic behaviour largely resembles other RNe in different phases. In the early phases of outburst broad emission lines dominated the spectra. The extensive wings reached FWZI$\\sim$8000-9000 km~s$^{-1}$ (Fig.\\ 1 and Fig.\\ 3). Later they became narrower. The low-resolution spectra covering the whole optical range show broad emission complexes $\\lambda\\lambda$4500--4700, $\\lambda\\lambda$4900--5100, $\\lambda\\lambda$5600--5800, $\\lambda\\lambda$5800--6000. Very similar spectra were obtained, e.g., for V394~CrA 5 and 6 days after maximum (Sekiguchi et al. 1989a, Williams et al. 1991, Williams 1992) and for U~Sco at +1.45 days (Anupama \\& Dewangan 2000) and between +3 and +5 days (Munari et al. 1999). Interestingly, Munari et al. (1999) observed exactly the same change of He II lines at 8237 \\AA\\ and 4686 \\AA\\ in U~Sco between +1.63 to +4.59 days (t$_2$=2.2 days, t$_3$=4.3 days) as has been found by us between +22 and +26 days (t$_2$=30 days, t$_3$=36 days). This fact suggests that the underlying physical mechanisms are similar.  The general appearance is typical for ``He/N'' spectra defined by Williams (1992). This gives further support to the statement of Williams (1992), that almost all novae with short recurrence times appear to have ``He/N'' type spectra (U~Sco, V394~CrA, V745~Sco, V3890~Sgr, LMC 1990 No. 2). The interpretation of this classification was given by Williams (1992) in terms of different components in the ejecta. The broader lines of the ``He/N'' spectrum originate in a discrete shell, which is ejected at considerably high velocities from the white dwarf surface at the peak of outburst.  In our dataset there are some weak pieces of evidence for another interpretation. The fairly strong and broad Fe II 5169 line around the maximum raises the possibility that CI~Aql belongs to the hybrid objects as discussed in Williams (1992). These stars have ``Fe II'' type spectra in the early phases that evolve to the ``He/N'' type. The very early ($\\Delta$t = $-$7 d) low-resolution spectrum taken by W. Liller showed no evidence of Fe II lines (Yamaoka et al. 2000), while our spectrum a week later showed strong and broad Fe II 5169 emission (Fig.\\ 1). Also, the intermediate resolution blue spectrum taken at +11 d (Fig.\\ 2) is very similar to what is observed in Nova LMC 1988 No. 2 by Sekiguchi et al. (1989b). Williams (1992) suggested that this behaviour is caused by two distinct phases of the outburst. The ``Fe II'' spectra probably originate from discrete and massive shells, which are optically thick, thus causing the early photosphere to occur in the ejected shell. As the expanding shell(s) becomes optically thin, the spectrum change to the ``He/N'' type with more rectangular line shapes. The double-structured absorption in the P-Cyg profiles (Figs.\\ 1 and 3) implies a quite complex inner structure of the ejecta, thus may strengthen the previous explanation.  The H$\\alpha$ evolution is also similar to what has been found in other RNe. Munari et al. (1999) presented a nice coverage between +0.64 d and +22.6 d for U~Sco, in which the early saddle-like H$\\alpha$ split into three components with velocity separation of the order of $\\pm$1600 km~s$^{-1}$. While the three components in the eclipsing system of U~Sco are difficult to interpret in terms of collimated beams of material ejected at a large angle from the orbital plane, the two components of CI~Aql can be explained with simple equatorial and polar rings of enhanced brightness in the ejected shell (Gill \\& O'Brien 1999). Since CI~Aql is an eclipsing system, we see the hypothetic rings nearly edge-on. As has been pointed out by Gill \\& O'Brien (1999), the high inclination means that the least information on the shell structure can be derived from the line profiles.  Finally, the main conclusions of this paper can be summarized as follows:  \\begin{itemize}  \\item Optical spectra taken between $-$0.6 and +53 d are presented and discussed. The overall appearance is similar to other recurrent novae, and consequently, the star belongs to the ``He/N'' type novae defined by Williams (1992). This means the lines originate from a discrete shell ejected by a velocity of 2000--2500 km~s$^{-1}$. Weak evidence for possible changing the type from ``Fe II'' to ``He/N'' is present in our dataset.  \\item We identified two diffuse interstellar bands which were used to estimate the interstellar reddening. The saturated interstellar Na I D doublet supports the fairly high colour excess adopted to be E(B$-$V)=0\\fm85$\\pm$0\\fm3.  \\item The visual light curve was used to determine epoch of maximum light and rates of decline. The recent light curve was compared with that of in 1917 and significant differences were found. t$_{\\rm 3}$ appears to be a more reliable indicator of the decline.  \\end{itemize}"
    },
    "0011/astro-ph0011156_arXiv.txt": {
        "abstract": "The physics of kinetic mixing between ordinary and mirror photons is discussed. An important role is played by four linear combinations we dub the physical photon, the sterile photon, the physical mirror photon, and the sterile mirror photon. Because of the mass degeneracy between the two gauge bosons, quantum coherence effects are important. The physical photon becomes a certain coherent superposition of the bare ordinary photon and the bare mirror photon. Similarly, the physical mirror photon is another, but {\\it not orthogonal}, coherent superposition. We discuss the physics of the interaction between physical mirror photons and ordinary matter. Observational signatures for some hybrid ordinary/mirror binary astrophysical systems are qualitatively discussed. We show that a small amount of ordinary matter at the center of a mirror star may make the mirror star observable. We speculate that the recently reported halo white dwarfs might actually be mirror halo stars. ",
        "introduction": "\\label{intro}  There has been considerable interest in the study of particle physics implications of the atmospheric \\cite{atmos}, solar\\cite{solar}, LSND \\cite{lsnd} and other neutrino experiments. The main purpose of such investigations has been to understand what theories could be responsible for the observed experimental features, which in the atmospheric and solar cases give strong evidence for large angle neutrino oscillations. One of these theories is the Exact Parity Model (EPM) (also known as the Mirror Matter Model) which introduces parity or ``mirror'' partners for all ordinary particles (except the graviton) and thus restores the parity invariance apparently broken by the weak interactions\\cite{ly,blin,flv}. The EPM predicts pairwise maximal mixing between ordinary and mirror neutrinos and provides a basis for an explanation of the atmospheric and solar neutrino  data\\cite{flv2}. This explanation does not require large mixing between generations and is also consistent with the LSND experiment.\\footnote{In the EPM the parity symmetry is unbroken by the vacuum, which is one of the reasons it is so simple and predictive. It is also possible to envisage models where the parity symmetry is spontaneously broken\\cite{spon} which typically have quite different features.}  Another interesting implication of this idea is that the mirror nucleons/atoms provide a natural candidate for dark matter\\cite{blin,hod,macho2}. This hypothesis can nicely explain some of the features of dark matter. For example, the heavy MACHO objects inferred to exist in the halo of our galaxy by the microlensing experiments\\cite{macho} can be naturally interpreted as mirror stars\\cite{macho2}. If the dark matter is composed of mirror matter then it is also possible that some mirror matter exists in our solar system and in other solar systems. In fact a mirror planet could form very close to an ordinary star and may provide a possible explanation for the close-in extra-solar planets discovered around nearby stars\\cite{plan}. For other possible astrophysical, cosmological and geophysical implications of mirror matter see Ref.\\cite{cos}.  Ordinary and mirror matter interacts through gravitation, and through the mixing of colourless and neutral particles with their mirror partners\\cite{flv,flv2,hol,gl}: neutrinos, the photon, the Z boson and the neutral Higgs boson can mix with their corresponding mirror states.\\footnote{ Other possibilities require physics beyond the minimal model. The most interesting of these would appear to be neutron -- mirror neutron and $K_L - K_L'$ mixing (the former requires baryon number violation, while the latter requires flavour changing neutral currents which could occur, for example, in a two Higgs doublet extension).} Of particular concern to this paper is photon -- mirror photon kinetic mixing, \\begin{equation} {\\cal L}_{\\text{int}} = - \\epsilon F_{\\mu\\nu} F'^{\\mu\\nu}, \\label{Lint} \\end{equation} where $F_{\\mu \\nu}$ ($F'^{\\mu\\nu}$) is the field strength tensor for ordinary (mirror) electromagnetism, and $\\epsilon$ is a free parameter. An important experimental consequence of Eq.(\\ref{Lint}) is the mixing of orthopositronium with mirror orthopositronium\\cite{gl}, leading to oscillations between these states in a vacuum experiment. The subsequent decays of the mirror state are invisible, resulting in an effective increase in the decay rate\\cite{gl}. A longstanding discrepancy between the theoretical prediction and some of the experimental measurements may in fact be resolved by this mirror world mechanism\\cite{fg} (see also Ref.\\cite{gen}). These experiments suggest the value $\\epsilon \\sim 5 \\times 10^{-7}$ which is too large to be palatable for standard BBN\\cite{cg}, but is nevertheless very interesting because of its terrestrially observable consequences.  The purpose of the present work is to consider in detail the physics of the photon -- mirror photon mixing. The degeneracy between the ordinary and mirror photons (both of which are assumed to be massless) leads to coherence effects which should be taken into account when considering various processes involving ordinary/mirror photon emission, absorption and scattering. Various interesting astrophysical implications will also be qualitatively discussed. In particular we will find that a small amount of ordinary matter at the center of a mirror star may make the mirror star observable. Indeed such objects may already have been observed, but interpreted as types of white dwarfs. ",
        "conclusions": "The physics of kinetic mixing between ordinary and mirror photons has been discussed. An important role is played by four linear combinations we have dubbed the physical photon, the sterile photon, the physical mirror photon, and the sterile mirror photon. Because of the mass degeneracy between the two gauge bosons, quantum coherence effects have to carefully considered. In particular, the physical photon becomes a certain coherent superposition of the bare ordinary photon and the bare mirror photon. Similarly, the physical mirror photon is another, but {\\it not orthogonal}, coherent superposition of the bare states. One consequence of this is that purely ordinary stars cannot lose energy through mirror photon emission.  We have discussed the propagation of mirror (ordinary) photons through an ordinary (mirror) matter medium. This led us to consider qualitative observational signatures of some hybrid ordinary/mirror binary astrophysical systems. In particular, a mirror star with an ordinary matter core can phenomenologically mimic a white dwarf, and we have speculated that badly understood objects such as DC white dwarfs, or the faint halo objects recently reported, might actually be such hybrids."
    },
    "0011/astro-ph0011010_arXiv.txt": {
        "abstract": "We explore in some detail the scenario proposed to explain the observed knee of the cosmic ray (CR) spectrum as due to the effects of photodisintegration of the CR nuclei by interactions with optical and soft UV photons in the source region. We show that the photon column densities needed to explain the experimental data are significantly lower than those obtained in previous estimations which neglected multinucleon emission in the photodisintegration process. We also treat more accurately the photodisintegration thresholds, we discuss the effects of photopion production processes and the neutron escape mechanism, identifying the physical processes responsible for the qualitative features of the results. This scenario would require the CR nuclei to traverse column densities of $\\sim 5 \\times 10^{27}- 2 \\times 10^{28}$~eV/cm$^2$ after being accelerated in order to reproduce the observed knee, and predicts that the CR composition should become lighter above $\\sim 10^{16}$~eV. ",
        "introduction": "It is widely accepted nowadays that cosmic rays (CRs) with energies per particle up to about $10^{18}$~eV are protons and nuclei of galactic origin. Furthermore, it is well established that the full CR energy spectrum has a power-law behavior with a steepening taking place at the so-called knee, corresponding to an energy $E_{knee}=3\\times 10^{15}$~eV. Although it is well known that the CR composition below the knee has a significant heavy component, its behavior beyond the knee remains somewhat controversial. Indeed, above $10^{14}$~eV/nucleus the information about the CR mass composition has to be drawn from extensive air shower observations at ground level with the results from different experiments not always being compatible. Moreover, the theoretical predictions depend on the hadronic model adopted and this can introduce further uncertainties on the inferred composition. The model dependence in certain observables can however be useful to discriminate among the hadronic models using the experimental data, and in this comparison QGSJET and VENUS turn indeed to be favoured against other models such as SIBYLL \\cite {fowler}. Thus, the CR mass composition beyond the knee is not definitively established yet, with some observations \\cite {swordy} suggesting that it turns lighter and others instead suggesting that the heavier components become dominant \\cite {amen1,ag00,ka01,ar00}.  The attempts made so far to explain the physical origin of the knee of the cosmic ray spectrum can be roughly classified into three kind of models. One of them exploits the possibility that the acceleration mechanisms could be less effective above the knee \\cite{fich, joki, koba}, while a second one assumes that leakage from the Galaxy plays the dominant role in suppressing CRs above the knee \\cite{syro,wdow,ptus}. The third scenario, originally proposed by Hillas \\cite{hillas}, considers nuclear photodisintegration processes (and proton energy losses by photomeson production), in the presence of a background of optical and soft UV photons in the source region, as the main responsible for the change in the CR spectrum. The first two scenarios are based on a rigidity dependent effect which consequently produces a change in the spectral slope of each nuclear component with charge $Z$ at an energy $\\sim E_{knee}Z$. Hence, it removes first the protons and the lighter nuclei and only at larger energies affects the heavier nuclei, predicting that the CR mass composition should become heavier beyond the knee. On the contrary, the third scenario predicts a lighter CR composition above the knee due to the disintegration of the propagating nuclei. The predictions of this last scenario have been worked out by Karakula and Tkaczyk \\cite{karak} some years ago. The aim of the present paper is to re-evaluate in detail the propagation of nuclei and protons in the source region following their approach, but taking into account some additional features that turn out to be important for the determination of the source parameters required for the scenario to work. In particular, we re-evaluate the nuclear photodisintegration rates taking into account multinucleon emission processes, we introduce more accurate threshold energies in the giant dipole resonance for every stable nucleus (according to recent remarks by Stecker and Salamon \\cite{steck1}), and discuss the impact of the neutron mechanism \\cite{neutron} which would allow the escape of neutrons from the source without further energy losses. Moreover, we show that by introducing a non negligible lower cutoff to the power-law photon distribution, the abrupt suppression of the all-particle CR spectrum above $10^{17}$~eV previously obtained \\cite{karak} is prevented, and hence the underlying reason for that effect is identified. Let us also mention that although we focus our study here on the scenario proposed to explain the observed steepening at the knee, we expect that our general considerations about the treatment of photodisintegration processes should also be relevant in other contexts. ",
        "conclusions": ""
    },
    "0011/astro-ph0011226_arXiv.txt": {
        "abstract": "Results from the third campaign of a ground-based multi-band observation of the millimeter emission of the sky from Tenerife (Canary Islands) are presented. The instrument consists of a 0.45 meter diameter off-axis telescope equipped with a 4-band multi-mode $^3He$ cooled photometer working at 1.1, 1.3, 2.1 and 3.1{\\notetoeditor{These are effective wavelengths}} mm wavelengths. The beam is well approximated by a Gaussian with 1$% ^{o}$.35 Full Width Half Maximum (FWHM) at all wavelengths. The wide wavelength coverage of our instrument allows us to characterize and reduce both the atmospheric and galactic contamination in our data. The CMBR data is analyzed in 6 multipole bands whose centers span the range $\\ell=39$ to $\\ell =134$ at the two longest wavelengths (2.1 and 3.1 mm). A likelihood analysis indicates that we have detected fluctuations in all bands at the two wavelengths. We have evidence of a rise in the angular power spectrum from low $\\ell $ to high $\\ell $. Our measured spectrum is consistent with current popular theories of large scale structure formation, COBE, and other recent balloon-borne experiments with similar wavelength coverage. ",
        "introduction": "Precise measurements of the angular power spectrum of the Cosmic Microwave Background Radiation (CMBR) provide strong constraints on cosmological parameters. After the era of the discovery of the anisotropy by COBE (\\cite{Smoot92} and \\cite{Bennet96}) we are now in the age of spatial spectroscopy of the primordial fluctuations. Many experiments \\notetoeditor{see $http://www.sns.ias.edu/% \\sim max/cmb/experiments.html$ and (\\cite{Page97})} have reported detections of anisotropy at different angular scales with evidence of excess of power at angular scales of about $\\ell \\sim 200$.  We designed an experiment sensitive to a broad angular range ($\\ell$ = 39 to 134 ). Our lowest $\\ell$-band $39^{+38}_{-24}$ partially overlaps with COBE's highest $\\ell$-band. The angular range also overlaps with COBE at our longest wavelength (3.1 mm). The partial overlap of the parameters of our experiment and COBE allows us to perform independent consistency checks.  In this letter, we present the data from our third observing campaign of mm-wave CMBR anisotropy from Tenerife (Canary Islands). The first campaign in 1993 was intended to test the instrumental set-up as well as characterize the site. The second campaign in 1994 (see {\\cite{Piccirillo97}} and {\\cite {fem97}}) measured the anisotropy at two $\\ell $ bands. In the third campaign we extended the angular scale coverage from 2 to 6 bands. Observations were carried out from May to July, 1996. ",
        "conclusions": "Improvements in our instrumental setup over our 1994 campaign, as well as refinements in our data analysis procedure, have allowed us to detect temperature fluctuations over a wide range of angular scales.  We have used data from our 4-band detector to reduce the atmospheric noise contamination and to delineate regions of our scans for which the galactic contamination is insignificant.   We find positive detections of residual temperature fluctuations in 6 angular ranges in our two longest wavelength bands (2.1 mm and 3.1 mm). These measurements, summarized in table 1, are our main result.   The derived temperatures  are consistent with those seen in other experiments . Our results essentially show a monotonically increasing amplitude of temperature fluctuation with l, and are consistent with some currently popular adiabatic theories of cosmological structure formation."
    },
    "0011/astro-ph0011540_arXiv.txt": {
        "abstract": "We present Giant Metrewave Radio Telescope HI 21cm absorption observations, of candidate and confirmed \\dla systems. The derived spin temperatures T$_{\\rm spin}$ are high, in all cases, $\\sim 1000$~K or higher. We have also collated from the published literature a list of damped absorbers for which 21cm observations exist, and discuss the implications of the observations for the nature of these systems.  A cross-comparison of the HI 21cm profiles (which trace the {\\it cold} gas) with the low ionization metal profiles (which trace {\\it all} the neutral HI, both cold and warm) shows that in all cases for which both spectra are available, the 21cm absorption coincides in velocity with the deepest metal line feature. This is consistent with models in which the deep metal line features arise from discrete clouds but not with models where the deepest features are the result of velocity crowding.  We also find that the typical derived spin temperatures of \\dla systems are considerably higher than those in the Milky Way or nearby spiral galaxies. The only exceptions are systems which are known to be associated with the disks of spirals; these {\\it do}, in fact, show low spin temperature. In a multi-phase medium the derived spin temperature is a weighted average of the temperatures of the individual phases. High apparent T$_{\\rm spin}$ values are hence to be expected from small, low metallicity objects since these objects should (as per existing theoretical models of the formation of a multi-phase ISM in the Milky Way and  high redshift proto-galaxies) have a lower fraction of the cold phase in their ISM as compared to large galaxies. The high  T$_{\\rm spin}$ is hence consistent with the observed low metallicities of \\dla absorbers as well as with recent findings that damped absorption is associated with a variety of galaxy types (as opposed to being confined to the disks of large spirals).  Finally, although the number of systems for which observations are available is small, we suggest that the following trend may be identified : at low redshift, \\dla absorption arises from a range of systems, including spiral galaxy disks, while, at high redshift, absorption occurs predominantly in smaller systems. ",
        "introduction": "\\label{sec:introduction} A gas cloud with a column density $N_{\\rm HI} > 10^{20}$ atoms/cm$^{-2}$ produces a broad Lyman$-\\alpha$ absorption profile with characteristic Lorentzian damping wings. Such systems can hence be easily identified from even modest resolution spectra of distant quasars. Although these high column density absorbers are extremely rare compared to the ubiquitous \\lya forest lines, they nonetheless make the dominant contribution to the observed {\\it neutral } hydrogen at high redshifts.  At $z=0$, the cross section for \\dla absorption is dominated by the disks of spiral galaxies (Rao \\& Briggs 1993). It is thus plausible that the damped systems seen at high redshift are either spiral galaxies or the precursors of spiral galaxies. In support of this hypothesis, large surveys of these objects have established that the mass density of neutral gas in \\dla systems at $z\\sim 3$ is comparable to the mass density of stars in luminous galaxies at $z\\sim 0$ (eg. Wolfe 1988, Lanzetta et al. 1991, Storrie-Lombardi et al. 1996). Further, the mass density in neutral gas today is far smaller than that at high redshift. The evolution of $\\Omega_{\\rm HI}$ with redshift thus matches the pattern expected from gas depletion due to star formation.  While the motivation for the original \\dla surveys was to find disk galaxies at high redshift (Wolfe et al. 1986), the nature of the systems in which this absorption arises remains as yet controversial.  Absorption line surveys principally establish the probability of finding an absorber along a random line of sight. This probability is proportional to the product of the volume number density of the absorbers and their typical cross-section; neither of the two is separately constrained. The observed frequency of occurrence of \\dla lines in QSO spectra at high $z$ is about 5 times larger than that expected from a population in which both the HI disk cross-section and the galaxy number density remain constant (Wolfe et al. 1986, Smith, Cohen \\& Bradley 1986, Lanzetta et al. 1991). The nominal conclusion is that, if damped absorption traces the precursors of spiral galaxies, then either the typical galaxy disk was somewhat larger in the past than it is now, or there were considerably more galaxies in the past than there are today.  Prochaska \\& Wolfe (1997,1998) attempted to settle the issue of the nature of damped absorbers in a purely observational manner, using high resolution spectra of low-ion metal transitions, to investigate the velocity structure of 31 systems. These low-ionization metals co-exist with neutral HI; their absorption profiles, however, are unsaturated and hence trace the kinematics of the neutral gas along the narrow line of sight to the background QSO. The authors found that the metal line profiles are systematically asymmetric (in that the deepest absorption preferentially occurs at the profile edge). Such asymmetries arise naturally in absorption from a spinning disk due to two effects (i) a ``radial effect'', where the strongest absorption occurs at the tangent point, i.e. due to velocity crowding, and (ii) a ``perpendicular effect'' in which the strongest absorption takes place in the midplane, if the line-of-sight intersects the midplane more than a few scale heights away from the major axis. The metal line profiles thus appear to support the hypothesis that \\dla absorption arises primarily in spiral galaxy disks. This is controversial, however, since the observed asymmetries can also be explained by merging proto-galactic clumps in hierarchical models (Haehnelt et al. 1998) and even by randomly moving clouds in a spherical halo (McDonald \\& Miralda-Escude 1999).  At low and intermediate redshifts, there have been a number of ground-based, as well as HST, observations to try and directly image the systems responsible for the absorption (see, for example, le Brun et al. 1997). From these observations, it appears that, contrary to predictions based on the HI content of $z=0$ galaxies, \\dla absorption is {\\it not} associated predominantly with bright ($L > L_{\\star}$) spiral galaxies, but appears to arise in galaxies spanning a wide range of morphological types. Interestingly, even in quasar-galaxy pairs at very low redshift, systems which show HI absorption appear to be predominantly disturbed or interacting galaxies (Carilli \\& van Gorkom 1992).  If the quasar against which the \\lya absorption is seen is radio loud, one can also attempt to detect redshifted 21cm absorption from the \\dla system (eg. Wolfe \\& Davis 1979, Wolfe, Briggs \\& Jauncey 1981, Wolfe et al. 1985, de Bruyn et al. 1996, Carilli et al. 1996, Lane et al. 1998, Chengalur \\& Kanekar 1999). Observing the  21cm line yields a number of interesting physical parameters. Firstly, since the 21cm line is not saturated, the absorption profile reflects the kinematics of the cold gas in the system. Further, (under suitable assumptions) the 21cm optical depth can be used along with the total HI column density (obtained from the \\dla profile) to yield the average spin temperature of the gas. Under a variety of astrophysical conditions (Field 1958), the spin temperature of the 21cm transition is identical to the kinetic temperature of the gas; it is hence a useful diagnostic. Finally, if the background source is extended, one can attempt to determine the transverse extent of the absorbing gas. This cannot be done using UV absorption lines since the background quasar is unresolved at these wavelengths.  We present, in this paper, Giant Metrewave Radio Telescope (GMRT) HI 21cm observations of candidate and confirmed low redshift \\dla systems. We have also collated from the published literature earlier 21cm studies of a set of such objects, and discuss the implications of these observations for the nature of the absorbers. ",
        "conclusions": "\\label{sec:dis}  We were able to collate from the published literature eighteen 21cm observations of \\dla absorbers, including cases where only upper limits have been obtained on the optical depth. This is probably {\\it not} an unbiased  sample since it is unlikely that all non-detections have been published. The relevant parameters of the sample are summarised in Table 2. Note that although 3C~446 and PKS~1451-375 are included in Table 2 and Figure 2, their status as bonafide \\dla systems remains to be confirmed.  The optical depth in the 21cm transition is a function of both the column density, $N_{\\rm HI}$, of the gas as well as its spin temperature, $T_s$. For an optically thin, homogenous cloud the exact relation is : \\begin{equation} \\label{eqn:Tspin} N_{\\rm HI} = { 1.823\\times10^{18} T_s \\over f} \\int \\tau_{21} dV \\end{equation} where $\\tau_{21}$ is the 21cm optical depth, $N_{\\rm HI}$ is the neutral hydrogen column density and $f$ is the covering factor; $T_s$ is in K and $dV$ in km s$^{-1}$. For an inhomogeneous absorber, the temperature $T_s$ in equation (\\ref{eqn:Tspin}) should be replaced by the column density weighted harmonic mean of the temperatures of the different components (assuming all components are optically thin). Thus, for example, if one has an absorber in which 50\\% of the gas is at 80~K and 50\\% at 8000~K, the measured spin temperature would be $\\sim 160$~K, while if 10\\% is at 80~K and 90\\% at 8000~K, the measured spin temperature would be $\\sim 730$~K.  In the case of damped systems, the column density $N_{\\rm HI}$ is known from the \\dla profile; this can be used in conjunction with the measured 21cm optical depth to obtain the spin temperature of the absorbing gas, if the covering factor $f$ is known. Many of the background sources listed in Table 2 are extremely compact making it likely that the covering factor is close to unity. Further, as discussed in more detail in Carilli et al. (1996), although the covering factor $f$ is sometimes poorly constrained, it is unlikely to lead to a systematic bias for the systems discussed here. We also emphasise that the sample of Table 2 does not contain unpublished non-detections of absorption; since $T_s$ and $\\tau_{21}$ are inversely related, the bias in the sample (if it exists) is thus towards {\\it lower} $T_s$ values. As we will see below, \\dla systems tend to have $T_s$ values higher than those observed in local spiral galaxies; the results from our sample are hence conservative since the bias tends to reduce this discrepancy.  \\subsection{Kinematics of damped systems} \\label{ssec:Kinematics}  A more fundamental consequence of the inverse relationship between $\\tau_{21}$ and $T_s$ is that 21cm absorption typically traces the cold gas ($T_s \\sim 100$ K) in the absorber. For example, in our galaxy, the warm neutral medium (WNM) has been detected in absorption only towards the very bright radio source Cygnus~A (Carilli et al. 1998), while, in extra-galactic systems, there has been only one claim of 21cm absorption that could arise from warm ($T_s \\sim 5000$~K) gas (Lane et al. 1999). Low ionization metal lines, on the other hand, trace both the warm and the cold neutral gas. A comparison between 21cm profiles and the profiles of metal lines can thus be used to glean further information on the phase structure and kinematics of the absorbing  system. Unfortunately, there are only four known 21cm absorbers for which high resolution ($\\sim$ 8 km s$^{-1}$) optical spectra are available. These are discussed individually below. \\\\  \\begin{table*} \\label{basicdata} \\caption{ Damped Ly-$\\alpha$ Systems with 21cm Observations} \\begin{tabular}{@{}lllllrll} Name & $z_{abs}$ & N$_{\\rm HI}$ & No. of & ${\\Delta V_{21}}^{  a}$ & T$_{\\rm s}^{  b}$ & Absorber ID& Ref\\\\ &&cm$^{-2}$& Components& km s$^{-1}$&K&& \\\\ &&&&&&& \\\\ 0738+313 (A)    &0.09123      &1.5e21    &3   &25          &825           & Dwarf ?         & 1,2 \\\\ 0738+313 (B)    &0.2212       &7.9e20    &1   &16          &1120          & Dwarf ?         & 1-3 \\\\ PKS~1413+135    &0.24671      &2.0e22    &1   &39          &150           & Spiral         & 4  \\\\ PKS~1451-375$^c$&0.2761       &1.3e20    &-   &-           &$>$ 1400      & -              & 5,6 \\\\ PKS~1127-145    &0.3127       &5.1e21    &7   &120         &910           & Group          & 3,6 \\\\ PKS~1229-021    &0.39498      &5.6e20    &2   &180         &170           & Spiral ?        & 7-10 \\\\ 3C196           &0.4366       &6.3e20    &3   &250         &-             & Barred Spiral  & 11-13 \\\\ 3C446$^c$       &0.4842       &6.3e20    &-   &-           &$>$ 1600      & -              & 5,6 \\\\ A0~0235+164     &0.524        &5.0e21    &5   &125         &170           & Spiral         & 14-17 \\\\ 3C286           &0.69215      &2.0e21    &1   &18          &1200          & LSB            & 9,18-20\\\\ PKS~0454+039    &0.8596       &5.0e20    &-   &-           &$>$870        & Galaxy         & 9,21,22 \\\\ MC3~1331+170    &1.77636      &7.6e21    &1   &25          &770           & -              & 23-25 \\\\ PKS~1157+014    &1.944        &6.3e21    &1   &60          &470           & Group          & 26-28 \\\\ PKS~0458-02     &2.03945      &8.0e21    &2   &30          &594           & -              & 29,30 \\\\ PKS~0528-2505   &2.8110       &2.2e21    &-   &-           &$>$ 730       & Group          & 31-33 \\\\ 2342+342        &2.9084       &2.0e21    &-   &-           &$>$ 1800      & -              & 31,34 \\\\ PKS~0336-017    &3.0619       &1.5e21    &-   &-           &$>$ 2500      & -              & 31,35\\\\ PKS~0201+113        &3.3875       &2.5e21    &-   &-           &$>$ 5600      & -              & 34,36-38 \\\\ &&&&&&& \\\\ \\end{tabular}  ${}^a$~Since many spectra have multiple-components and are not well fit by a single Gaussian, the width quoted here is the entire velocity range over which absorption is seen. ${}^b$~The spin temperature (or its limits) quoted here has been consistently (re)computed from the published data using equation (\\ref{eqn:Tspin}), and the latest available values for all parameters in the equation. In certain instances, this value may differ slightly from that quoted by the authors in the original reference. The only exception is PKS~0201+113, for which multiple and quite different values are quoted in the literature. We use the value from Kanekar \\& Chengalur (1997). ${}^c$~Candidate \\dla systems, based on IUE spectra.   ${}^{1}$~Rao \\& Turnshek (1998), ${}^{2}$~Chengalur \\& Kanekar (1999), ${}^{3}$~Lane et al. (1998), ${}^{4}$~Carilli et al. (1992), ${}^{5}$~Lanzetta et al. (1995), ${}^{6}$~This paper, ${}^{7}$~Brown \\& Spencer (1979), ${}^{8}$~Briggs (1998), ${}^{9}$~Boiss\\'{e} et al. (1998), ${}^{10}$~le Brun et al. (1997), ${}^{11}$~Brown \\& Mitchell (1983), ${}^{12}$~de Bruyn et al. (2000), ${}^{13}$~Cohen et al. (1996), ${}^{14}$~Roberts et al. (1976), ${}^{15}$~Wolfe et al. 1982, ${}^{16}$~Burbidge et al. (1996), ${}^{17}$~Cohen et al. (1999), ${}^{18}$~Brown \\& Roberts (1973), ${}^{19}$~Davis \\& May (1978), ${}^{20}$~Steidel et al. (1994), ${}^{21}$~Steidel et al. (1995), ${}^{22}$~Briggs \\& Wolfe (1983), ${}^{23}$~Wolfe \\& Davis (1979), ${}^{24}$~Wolfe et al. (1976). ${}^{25}$~Carswell et al. 1975 ${}^{26}$~Wolfe et al. (1981), ${}^{27}$~Briggs et al. (1984), ${}^{28}$~Fynbo et al. (1999), ${}^{29}$~Wolfe et al. (1985), ${}^{30}$~Briggs et al. (1989), ${}^{31}$~Carilli et al. (1996), ${}^{32}$~M{\\o}ller \\& Warren (1993), ${}^{33}$~M{\\o}ller \\& Warren (1998), ${}^{34}$~White et al. (1993), ${}^{35}$~Lu et al. (1993), ${}^{36}$~de Bruyn et al. (1996), ${}^{37}$~Briggs et al. (1997), ${}^{38}$~Kanekar \\& Chengalur (1997). \\end{table*}  \\noindent 1. 1229-021 ($z_{abs} = 0.3950$) The metal lines of this system were studied by  Lanzetta \\& Bowen (1992), using the AAT with a velocity resolution of $7$~km s$^{-1}$. The 21cm spectrum (Brown \\& Spencer 1979; Briggs 1998) shows a narrow component centered at the redshift of the deepest part of the metal line profile. The metal lines show two more components at $+100$~km s$^{-1}$ and $+200$~km s$^{-1}$, relative to $z = 0.395$. While the 21cm profile has a broad shallow component with a total width between nulls of $180$~km s$^{-1}$ there is no absorption associated with the $+200$~km s$^{-1}$ metal line component nor a separate discrete component at $+100$~km s$^{-1}$.  \\noindent 2. 3C286 ($z_{abs}$ = 0.69215) The velocity structure of the damped system at $z \\sim 0.692$ towards 3C286 is exceedingly simple in both the 21cm and low-ion (FeII, CaI) transitions. Each consist of a single narrow component with b values of $5.2$~km/s for the 21cm line (Davis \\& May 1978) and $6.5$~km/s for the FeII line (Meyer \\& York 1992). The metal lines occur at $z = 0.69218 \\pm 0.00007$ while the 21cm line is at $z = 0.692154 \\pm 0.000002$; thus, the locations are coincident in velocity, within the error bars.  \\noindent 3. 1331+170 ($z_{abs} = 1.7764$) The SiII 1808 line detected from this system (resolution $\\sim 6.6$~km/s) is clearly asymmetric (Prochaska \\& Wolfe 1997), with its deepest feature at $z = 1.77636 \\pm 0.00006$ . The 21cm profile (Wolfe \\& Davis, 1979) has a single component at $z = 1.77642 \\pm 0.0001$, the deepest part of the SiII line coincides with the 21cm line, again within the error bars. However, once again there are discrete components in the SiII spectrum which have no counterparts in 21cm absorption.  \\noindent 4. 0458-020 ($z_{abs} = 2.039$) The 21cm line (Wolfe et al. 1985) has two clear components, at $z = 2.03937$ and $z = 2.03955$. The two deepest features of the CrII~2056 profile (Prochaska \\& Wolfe 1997) coincide with the 21cm features (within the error bars); however, the deeper CrII feature corresponds to the weaker 21cm feature (see Figure 2, in Prochaska (1999), for an overlay of the two profiles). There are also several weaker CrII features which have no corresponding features in the 21cm spectrum.  Four systems are clearly too small a sample to draw any far-reaching conclusions regarding the general kinematics of damped absorbers. It is, however, revealing that the strongest metal line absorption coincides in velocity with a 21cm absorbing component, in every case for which high-resolution optical spectra are available. This points to a physical coincidence between the cold gas responsible for the 21cm absorption and the (not necessarily cold) gas giving rise to the deepest features in the low-ion profiles. This is consistent with any model for which the absorption arises in discrete clouds, but not with models for which the deep features in the low-ion profiles are the result of velocity crowding.  The other noteworthy feature is that the low ionization metal lines are in general much broader than the 21cm profiles, and often contain discrete velocity components which have no analogues in the 21cm spectrum. For example, if one assumes that the neutral hydrogen column density scales with the depth of the low ionization lines (which is {\\it not} necessarily a good assumption), the noise levels on the 21cm spectrum towards 0458-020 (Wolfe et al. 1985) are sufficiently low to detect 21cm absorption corresponding to the feature at $v \\sim 50$ km s$^{-1}$ in the CrII spectrum of Prochaska \\& Wolfe (1997) (see Figure 2 of Prochaska (1999)). Much as in our own ISM, the components with no corresponding 21cm absorption probably arise in warm gas which contributes to the total $N_{\\rm HI}$ measured in the Ly$-\\alpha$ line, but whose spin temperature is too high for it to be seen in 21cm absorption. It is possible, of course, that the above UV absorption arises from a number of small cold clouds, with high inter-cloud velocity dispersion, with the column density of each individual cloud being too low to be detected in the 21cm line. This is, however, an unlikely scenario since the clouds would require a column density of $\\sim 10^{19}$ cm$^{-2}$ for self-shielding and column densities as low as $\\sim 10^{18}$ of {\\it cold} HI can typically be detected in 21cm absorption.  \\begin{figure} \\centering \\psfig{file=tvsz.ps,width=3.5truein,height=3.0truein,angle=-90} \\caption{ Plot of $T_s$ for all damped systems for which published 21cm absorption measurements exist. Upper limits are shown by upward pointing arrows. Absorbers identified as spiral galaxies are shown as hollow squares.} \\label{fig:Tspin} \\end{figure}  \\subsection{The spin temperature of damped absorbers} \\label{subsec:Tspin}  Figure~\\ref{fig:Tspin} shows a plot of spin temperature v/s redshift for 17 systems from Table 2 (excluding 3C196, for which the background source is extended, making the column density uncertain). It can be clearly seen that the majority of systems tend to have $T_s$ values much larger than 500~K. For comparison, spin temperatures in the disc of the Galaxy or Andromeda are typically between 100 and 200 K (Braun \\& Walterbos 1992; Braun 1997). Similarly, Carilli \\& van Gorkom (1992) find that when the line-of-sight passes through the disc of a galaxy, the spin temperatures derived from 21cm emission-absorption arising in nearby quasar-galaxy pairs are also low, i.e. within 1 $\\sigma$ of 300 K, (although the error bars are admittedly large).  Further, as discussed by Carilli et al.~(1996), the difference in $T_s$ values between damped systems and local discs is probably not due to the different methods used for the measurements in the two cases.  Interestingly enough, however, {\\it all} the absorbers in Table~2  which have been identified as spiral galaxies either have low $(<$ 200 K) spin temperatures or, where the column density is not independently known (in the case of the extended source, 3C196), have absorption profiles consistent with low $T_s$ values for moderate values of column density (de Bruyn, Briggs \\& Vermeulen 2000). It thus appears that low spin temperatures {\\it do} occur in damped Lyman-$\\alpha$ systems, when the line-of-sight passes through the disk of a spiral galaxy. This indicates that the discrepancy in spin temperature between the majority of damped absorbers and local spirals is a real one, which may well arise from a true difference in the physical characteristics of the two classes of systems. The likely explanation for the high spin temperatures in damped systems is (as suggested earlier, eg. Carilli et al. 1996) that their fractional content of warm gas is larger than that of local spirals.  In the Milky Way, neutral gas can exist in pressure equilibrium in two stable phases, the warm neutral medium (WNM) at T~$\\sim 10^4$~K and the cold neutral medium (CNM) at T $\\sim 60$~K. At high pressures (such as those obtained in the midplane  of the Galaxy), only the cold phase exists, while, at low pressures (as obtained at  large scale heights, far away from the midplane), only the warm phase is found.  At intermediate pressures, the two phases co-exist in pressure equilibrium. Detailed models of heating and cooling processes in a two-phase neutral medium (Wolfire et al. 1995) have shown that the pressure required to produce the  cold phase increases with a decrease in metallicity. The cold phase should hence  be rare in smaller galaxies which have both lower midplane pressures and lower metallicities. Since the fractional contribution of the CNM to the total column density would be smaller in such systems, they would (as detailed earlier) have higher average $T_s$ values than spiral galaxies.  The metallicity evolution of \\dla systems is currently controversial. Damped absorbers typically have low metallicity ($\\sim 0.1$ solar) and there is little direct evidence for evolution of metallicity with redshift (Pettini et al. 1999). However Savaglio \\& Panagia (1999) find that if one assumes, and fits for, depletion using depletion patterns observed in the Milky Way, the metallicity does appear to increase with decreasing redshift. If the absorbers were indeed spirals, one would have expected to see metallicity evolution. If, on the other hand, damped systems typically arise in smaller galaxies, a consistent picture can be seen to emerge. Smaller galaxies have lower levels of star formation and are hence generally metal poor. Further, their midplane pressures are lower than those of large galaxies, making it doubly difficult to create the conditions for the formation of the CNM. Thus, both the low metallicity and the high spin temperature of damped systems can be consistently explained if damped Lyman-$\\alpha$ absorption occurs predominantly in small rather than large galaxies.  Finally, a detailed analysis of the formation of a multi-phase medium in proto-galaxies (Spaans \\& Norman 1997, Spaans \\& Carollo 1997)  indicates that the epoch of formation of a multi-phase ISM is a strong function of the galactic mass; while galaxies with baryonic mass $\\sim 5 \\times 10^{11}$~M$_\\odot$ form the cold phase of the ISM by redshifts of $\\sim 3$, dwarf galaxies with mass $\\sim 3\\times 10^9$~M$_\\odot$ typically form the cold phase only by redshfits $\\sim 1$.  It can be seen from Figure~\\ref{fig:Tspin} that low spin temperatures are not obtained at high redshift ($z \\ga 2$), while both low and high $T_s$ values are found at low redshift. Further, Table~\\ref{basicdata} indicates that large velocity widths ($\\Delta V_{21} > 100$ km s$^{-1}$) are not found for $z \\ga 2$, while both large and small widths are seen at low redshift. This is qualitatively consistent with what is expected in hierarchical merging models, (eg. Kauffmann 1996). Interestingly, $1127-145$ is the only system which has both a large 21cm velocity spread and a high temperature; all other systems with  large velocity widths are identified with low $T_s$ spirals. There are, however, at least three galaxies at the redshift of the \\dla absorber (Lane et al. 1998). It is possible that they are interacting and that some of the HI in this system is associated with tidal features and has a high spin temperature.  Hence, although the numbers are small, the general trends in $T_s$ and $\\Delta V_{21}$ are consistent with scenarious in which \\dla absorbers are typically small systems at high redshifts, while at low redshifts, they are a composite population including both spiral galaxies as well as smaller systems. This would be consistent both with hierarchical models of galaxy formation and narrow band imaging observations (Fynbo et al. 1999), which show that high redshift \\dla systems often have several Ly-$\\alpha$ emitting objects at the redshift of the absorber.  \\noindent {\\bf Acknowledgments} The observations presented in this paper would not have been possible without the many years of dedicated effort put in by the GMRT staff to build the telescope. We are also grateful to the referee for useful comments."
    },
    "0011/astro-ph0011406_arXiv.txt": {
        "abstract": "We present a new scenario for the behaviour of Be/X-ray binaries based on long-term multiwavelength monitoring and the decretion disc model. The circumstellar discs of the primaries are truncated because of the tidal and resonant effect of the neutron star. The geometry of the systems and the value of viscosity determine the presence or absence of Type I X-ray outbursts. The interaction of a strongly disturbed disc with the neutron star originates Type II X-ray and optical outbursts. ",
        "introduction": "Be/X-ray binaries are composed of a neutron star orbiting a Be star and accreting from its circumstellar disc. The high-energy radiation is believed to arise due to accretion of material associated with the Be star by the compact object (see Negueruela 1998; see also Bildsten et al. 1997). Some Be/X-ray binaries are persistent X-ray sources (see Reig \\& Roche 1999), displaying low-luminosity ($L_{{\\rm x}} \\sim 10^{34}\\:{\\rm erg}\\,{\\rm s}^{-1}$) at a relatively constant level (varying by up to a factor of $\\sim 10$). On the other hand, most known Be/X-ray binaries (though this is probably a selection effect) undergo periods in which the X-ray luminosity suddenly increases by a factor $\\ga 10$ and are termed Be/X-ray transients. The distinction between the two groups is difficult to establish solely on terms of the temporal behaviour, since some sources such as A\\,1118$-$61 and 4U\\,1145$-$619 display relatively weak outbursts, though there is some evidence that the X-ray spectral properties of the two groups could be different.  Be/X-ray transients fall along a relatively narrow area in the $P_{{\\rm orb}}$/$P_{{\\rm spin}}$ diagramme (see Corbet 1986; Waters \\& van Kerkwijk 1989), indicating that some mechanism must be responsible for the correlation. Those systems with fast-spinning neutron stars do not show pulsed X-ray emission during quiescence (though non-pulsed radiation could be caused by accretion on to the magnetosphere) because of the centrifugal inhibition of accretion (Stella et al. 1986). Systems with more slowly rotating pulsars show X-ray emission at a level $L_{{\\rm x}} \\la 10^{35}$ erg s$^{-1}$ when in quiescence. Transients show two different kinds of outbursts:  \\begin{itemize} \\item Moderate intensity X-ray outbursts ($L_{{\\rm x}} \\approx 10^{36} - 10^{37}$ erg s$^{-1}$) occurring in series separated by the orbital period (Type I or normal), generally (but not always) close to the time of periastron passage of the neutron star.  The duration of these outbursts seems to be related to the orbital period. \\item Giant (or Type II) X-ray outbursts ($L_{{\\rm x}} \\ga 10^{37}$ erg s$^{-1}$), lasting several weeks. The parameters of these outbursts do not correlate clearly with orbital parameters, though in A\\,0535+26 and 4U\\,0115+63 they seem to start always a few days after periastron passage. \\end{itemize} ",
        "conclusions": ""
    },
    "0011/astro-ph0011138_arXiv.txt": {
        "abstract": "{ We discuss here the pulsation properties of the $\\delta$ Scuti star BV Circini on the basis of data obtained during a simultaneous photometric and spectroscopic campaign in 1996 and a spectroscopic one in 1998, and taking also advantage of the previous photometric observations by Kurtz (1981). Nine pulsation modes were detected from photometry and thirteen from spectroscopy; five of them are in common to both techniques. The spectroscopic data give ample evidence of dramatic amplitude variations in some modes, in particular the strongest spectroscopic mode in 1998 was not detectable in 1996 data. The two dominant photometric modes (6.33 and 7.89 \\cd) are observed on both seasons. The typing of the modes was performed by means of a simultaneous model fit of line profile and light variations. The 6.33 \\cds photometric term is probably the fundamental radial mode, while the 7.89 \\cds is a nonradial mode with $m\\neq$0. There are six high--degree prograde modes with an azimuthal order $m$ ranging from --12 to --14, and also a retrograde mode with $m\\sim 7$. These modes combined with the identification of the 6.33 \\cds mode allowed us to estimate $i\\sim 60^o$ for the value of the inclination of the rotation axis. An accurate evaluation of the main stellar physical parameters is also proposed as a result of the pulsational analysis. \\keywords {Methods: data analysis -- techniques: spectroscopic -- techniques: photometric -- stars: individual: BV Cir -- stars: oscillations -- stars: variables: $\\delta$ Sct } } \\authorrunning{Mantegazza et al.} \\titlerunning{BV Cir} ",
        "introduction": "Asteroseismology is knowing an increasing interest thanks to satellites which are expected to fly in the incoming years (MOST, MONS, COROT) and which will perform accurate investigations of stellar variability at the $\\mu$mag level. Stars located in the lower part of the instability strip are suitable candidates to find a large quantity of excited modes, both radial and nonradial. The detailed knowledge of the pulsational properties of $\\delta$ Sct stars can help in the preparation of the scientific background of these missions; a review of the photometric properties of these variable stars is given by Poretti (2000). The main difficulty in their investigation is the typing of the excited modes, i.e. classifying the oscillations in terms of quantum numbers $n, \\ell$ and $m$. To proceed on this way, we started an observational project combining photometric and spectroscopic techniques at the European Southern Observatory. In this paper we supply the last result, after those on FG Vir (Mantegazza et al. 1994), X Cae (Mantegazza \\& Poretti 1996, Mantegazza et al. 2000) and HD~2724 (Mantegazza \\& Poretti 1999).  BV Circini$\\equiv$ HD 132209 was discovered to be a $\\delta$ Scuti variable star by Kurtz (1981). He detected four pulsation modes from the frequency analysis of 18 nights of observations in the $B$ and $V$ colours in 1980. The author also pointed out that these modes did not provide a complete description of the stellar variability. The strongest term at 6.328 \\cd\\, was suggested to be the radial fundamental mode. Since the star is rather bright ($V=6.56$) and its light variations are quite large ($\\Delta V>0.1$ mag), we decided to include it among our targets. ",
        "conclusions": "BV Circini has a lot of excited pulsation modes with observed frequencies ranging from low to high values. Nine frequencies were detected from photometric data (1980 and 1996) and 13 from spectroscopic (1996 and 1998) ones, 5 of them are in common in both datasets. The amplitudes of some modes are dramatically changing and this is particularly evident from the two spectroscopic datasets, in particular the strongest spectroscopic mode in the 1998 data at 17.28 \\cds, was not detectable in the previous dataset.  Six of the detected modes are high azimuthal order prograde modes with $-14\\leq m \\leq -12$ and  one is retrograde with $m=$7 or 8. For only two of the modes detected by simultaneous photometry and spectroscopy it has been possible to suggest a complete $(\\ell, m)$ identification, while better $S/N$ spectroscopic data are needed for the others. The two identified modes are the one at 6.33 \\cds (the photometric dominant mode), which can be  identified as the  $\\ell=0$ (radial mode). The other is the one at 7.89 \\cds which is a low degree prograde mode with $\\ell=1$ or $2$ and $m=-1$. The proposed typings of the pulsation modes of BV Cir are summarized in Tab.~\\ref{summa}.  By using the physical parameters computed in Sect.~2 we can evaluate the pulsation constant of the 6.33~\\cds mode, i.e. $Q=0.032\\pm0.005$~d. Therefore, this mode is probably the radial fundamental in agreement with the suggestion by Kurtz (1981). If we take this result for granted, we can use this information to narrow the uncertainties on some physical parameters. In fact, assuming for this mode the theoretical value of $Q=0.033\\pm 0.001$~d as pulsation constant, we get from its definition $\\overline{\\rho}/\\overline{\\rho}_{\\sun}=0.044\\pm 0.003$ and also $\\log g=3.63\\pm0.04$. Moreover if we combine the mean density with the radius estimate we get $M=2.04\\pm 0.12 M_{\\sun}$, again in excellent agreement with the value derived from the evolutionary models (see Sect.~2). \\begin{table} \\caption[]{Basic physical parameters of BV Cir} \\begin{flushleft} \\begin{tabular}{ll} \\hline\\noalign{\\smallskip} Parameter & Estimated value \\\\ \\hline Mass & $2.04\\pm 0.12 M_{\\sun}$\\\\ Effective temperature & $7260\\pm100^oK$\\\\ Absolute bolometric magnitude & $0.97\\pm 0.17$\\\\ $\\log g$& $3.63\\pm0.04$\\\\ Radius &$3.59\\pm0.30 R_{\\sun}$\\\\ Equatorial rotational velocity & $111\\pm 12$ \\kms\\\\ Rotational frequency & $0.61\\pm0.07$ d$^{-1}$\\\\ \\hline \\end{tabular} \\label{para} \\end{flushleft} \\end{table}  Hence by comparing the information from pulsation properties (as derived from our analysis), {\\sc hipparcos} parallax and effective temperature (as derived from multicolour photometry), we determined  physical parameters which are perfectly consistent (Tab.~\\ref{para}.)  We have remarked in the previous section that the discriminants of the detected modes are not particularly sensitive to the inclination of the rotation axis. However we can constrain this value using the azimuthal orders derived for the high degree modes and the fact that the 6.33 \\cd mode is the radial fundamental, because in this case their frequencies in the stellar reference frame should be larger than the radial fundamental one. Among the prograde modes the critical one is the 13.85~\\cds, which with $m=-12.5\\pm 1$ imposes $i\\geq 62^o\\pm10^o$, while the 1.48~\\cds with $m=7.5\\pm1$ gives $i\\leq 56^o\\pm 12^o$. By combining the two results we obtain $i\\sim 60^o$. Consequently the equatorial rotation velocity is about 111 kms$^{-1}$ and the rotation frequency $\\Omega$ about 0.61~rev~d$^{-1}$.  Let us finally deal with a few considerations about the technique to fit line profile variations. We have seen that in order to get an unambiguous discrimination between the possible identifications we need data with a rather high $S/N$. The present data, with a $S/N$ of the order of 250, supply useful information for the strongest spectroscopic terms only. Since the observed amplitudes are typical of $\\delta$ Scuti stars, it is not unreasonable to believe that, in order to do a good job, data with a $S/N$ of 500 or better could be necessary.  We have also explored the difference between fitting line profile variations considering velocity fields only or taking also into account flux variations. We have seen that in the first case we tend to overweight the higher degree modes when we try to identify low degree modes. For these modes it is therefore important to consider flux variations too, and in this case, in order to properly constrain the fitted model, it is necessary to have simultaneous light variations."
    },
    "0011/astro-ph0011412_arXiv.txt": {
        "abstract": "We have mapped the total sky brightness at \\fb\\ in two adjacent $60\\deg$ declination bands with the portable 5.5-m parabolic reflector of the Galactic Emission Mapping (GEM) project, an on-going international collaboration to survey the radio continuum of the sky in decimeter and centimeter wavelengths. The observations were conducted from two locations, one in the USA and the other in Brazil, using a novel instrumental approach to overcome the well-known shortcomings of survey experiments. Our strategy consists of a 1-rpm rotating dish to circularly scan the sky at $30\\deg$ from zenith. The dish uses a rim-halo to re-direct the spillover sidelobes of its backfire helical feed toward the sky and the entire assembly has been enclosed inside a wire mesh ground shield in order to minimize and level out the contamination from the ground. The diffraction characteristics of this set-up have been succesfully modelled and undesired systematic striping across the observed bands has been carefully removed by a baseline propagation method which exploits the time-forward and time-backward intersections of the circular scans. The map displays nearly 300~hours of our best quality data taken with a HPBW of $5.4\\deg$ at a sensitivity of 20 mK. ",
        "introduction": "Today, as we witness the dawn of a precision era in observational cosmology, we are becoming increasingly aware that our present knowledge and understanding of the Galactic foregrounds demand a more thorough underpinning. During the course of this meeting we have learned about the results from the latest experiments for measuring the anisotropy in the Cosmic Microwave Background Radiation (CMBR). Yet, despite their success in ruling out a significant contamination of Galactic origin, there are still opened questions regarding the spatial and spectral distribution of the Galactic constituents. New insights have already improved on this picture as the work of Lazarian on Galactic dust properties and of Reynolds on Galactic free-free mapping through H$\\alpha$ association have been able to tell us. Such an effort also lies at the core of the Galactic Emission Mapping (GEM) project in order to reveal a new look of the Galactic synchrotron component, capable of overcoming the well-known shortcomings of the existing surveys; let alone the important implications for the search of spectral distortions in the CMBR. ",
        "conclusions": ""
    },
    "0011/astro-ph0011248_arXiv.txt": {
        "abstract": "We present a new model for the Galactic population of close double white dwarfs. The model accounts for the suggestion of the avoidance of a substantial spiral-in during mass transfer between a giant and a main-sequence star  of comparable mass and for detailed cooling models. It agrees  well with the observations of the local sample of white dwarfs  if the initial binary fraction is $\\sim\\!50$\\,\\% and an {\\it ad hoc} assumption is made that white dwarfs with $M\\!\\aplt 0.3\\,\\ms$ cool faster than the models suggest. About 1000 white dwarfs  with $V\\!\\aplt\\!15$ have to be surveyed for detection of a pair which  has $M_1+M_2\\!\\apgt\\!\\mc$ and will merge within 10\\,Gyr. ",
        "introduction": "The interest in the close double white dwarfs (hereafter CDWD) stems from several reasons: (i) white dwarfs are the endpoints of stellar evolution; (ii)  CDWD experienced at least two stages of mass exchange and thus provide an important tool for testing the evolution of binaries; (iii) merging double CO white dwarfs are a model for SNe\\,Ia; (iv) CDWD may be the most important contributors to the gravitational wave noise  at $\\nu \\aplt 0.01\\,{\\rm Hz}$, possibly burying  signals of other sources.  We review the data on the currently known CDWD and present results of the modelling of the population of white dwarfs, which involves several new aspects. The most important are the treatment of the first stage of mass loss without significant spiral-in \\cite{nvy+00}, the use of detailed models for the cooling of white dwarfs and the consideration of different star formation histories for the Galactic disc. ",
        "conclusions": "1. The standard Algol-type evolution and spiral-in in common envelopes  cannot explain the wide orbits of the progenitors of close double WD  after the first stage of mass loss. In  low-mass binaries with similar masses of components,  the loss of the envelope of the initial primary probably doesn't cause a strong spiral-in. In the absence of a physical picture of this process, we suggest to describe the change of separation of the components by an angular momentum loss law [Eq.~(\\ref{eq:aml})]. An attempt to model the parameters of three observed CDWD with two helium components results in $1.4 \\aplt \\gamma \\aplt 1.8$.  2.  Modelling of the observed sample of close binary WD shows that low-mass WD have to cool faster than is suggested by the recently published cooling curves for helium WD with thick hydrogen envelopes \\cite{dsb+98,sea00}. We suggest that an additional mass loss by WDs in winds or thermal flashes may strongly diminish the masses of envelopes and reduce the cooling times.  3. A reasonable agreement of the population synthesis results with observations in  the estimates of the local space density of WD, and the  $P_{orb}- m$, $P_{orb} -q$, and $N(P_{orb})$ distributions for  CDWD may be achieved if the parameter $\\gamma$ in Eq.~(\\ref{eq:aml})  is  $\\sim\\!1.7$ and it's assumed that WD with $M\\!\\leq\\!0.3$\\,\\ms\\ cool like the most massive helium WD.  Further reduction of cooling times for the helium WD may even improve agreement.  4. The model with  an initial binary fraction of 50\\%\\ (2/3 of stars are in binaries) fits the observations better than  the model in which  initially 100\\% of stars are in  binaries.   5. We estimate that the detection  of at least one WD pair  with $M_1+M_2 \\geq M_{Ch}$ may require  a survey of $\\sim 1000$ white dwarfs. But  the uncertainties in the input data for the population synthesis studies and badly known selection effects, most probably make this estimate accurate only within a factor of $\\sim 2\\div 3$.  \\vspace{0.3cm} LRY  acknowledges the support of the LOC and Astronomical Institute ``Anton Pannekoek'' which enabled him to attend the conference.  This work was supported by NWO Spinoza grant 08-0 to E.P.J.~van den Heuvel, RFBR grant 99-02-16037, ``Astronomy and Space Research'' Program (project 1.4.4.1), and NASA through Hubble Fellowship grant HF-01112.01-98A awarded (to SPZ) by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under contract NAS\\, 5-26555."
    },
    "0011/astro-ph0011281_arXiv.txt": {
        "abstract": "Results are presented of a new VLA$-$ROSAT study that probes the magnetic field strength and distribution over a sample of 16 ``normal'' low redshift ($z \\leq 0.1$) galaxy clusters. The clusters span two orders of magnitude in X-ray luminosity, and were selected to be free of (unusual) strong radio cluster halos, and widespread cooling flows. Consistent with these criteria, most clusters show a relaxed X-ray morphology and little or no evidence for recent merger activity.  Analysis of the rotation measure (RM) data shows cluster-generated Faraday RM excess out to ${\\sim} 0.5\\, h_{75}^{-1}$ Mpc from cluster centers. The results, combined with RM imaging of cluster-embedded sources and ROSAT X-ray profiles indicates that the hot intergalactic gas within these ``normal'' clusters is permeated with a high filling factor by magnetic fields at levels of $<|B|>_{icm}$ = 5$-$10 $(\\ell /10\\ {\\rm kpc})^{-1/2}\\,h_{75}^{1/2}\\ \\mu $G, where $\\ell$ is the field correlation length. These results lead to a global estimate of the total magnetic energy in clusters, and give new insight into the ultimate energy origin, which is likely gravitational. These results also shed some light on the cluster evolutionary conditions that existed at the onset of cooling flows. ",
        "introduction": "Determinations of magnetic field strengths in the intracluster medium over the past decade have revealed fields of unanticipated strength in some clusters. Additionally, these studies have raised apparent contradictions between field strengths (or limits) obtained from different observational methods. In order to understand the physical conditions in the intracluster medium, it is important to understand the origin(s) of cluster fields, and the astrophysical processes that relate them to other energy constituents of the cluster gas.  Estimates of magnetic field strengths in the intracluster medium (ICM) can be determined from measurements of Faraday rotation through the intracluster gas, and an independent measurement of the ICM thermal electron density, $n_{e}$. Faraday rotation is given by: \\begin{equation} {\\rm RM}=\\frac{\\Delta \\chi}{\\Delta \\lambda^2}=811.9\\int_0^L n_{e} B_{\\|}d\\ell \\ \\ {\\rm rad\\ m^{-2}}, \\label{eqn:rm} \\end{equation} where $\\chi$ is the position angle of the linearly polarized radiation at wavelength $\\lambda$, $n_{e}$ is the thermal electron density in ${\\rm cm^{-3}}$, $B_{\\|}$ is the line of sight magnetic field strength in $\\mu$G, and $L$ is the path length in kpc. The thermal electron density in a cluster can be determined from X-ray surface brightness profiles of the hot (T${\\rm \\sim 10^8}$ K), diffuse (${n_{e}}\\sim 10^{-3}\\,h_{75}^{1/2}\\ {\\rm cm^{-3}}$) gas \\citep{hans95} which fills the cluster potential.  The first study of background RMs over a {\\it single} cluster \\citep{kk90} was a targeted set of deep VLA observations of 18 sources close in angular position to the Coma cluster (which has no cooling flow, but a strong synchrotron halo). The result of the study was $ <|B|>_{icm}$ = 2.5 $(\\ell /10\\ {\\rm kpc})^{-1/2}\\,h_{75}^{1/2}\\ \\mu $G where $\\ell$ is the B-reversal scale. In a subsequent study by \\citet{feretti95} the discovery of smaller $\\ell$ scales down to 1 kpc raised this estimate to $\\sim 7.2\\,h_{75}^{1/2}\\mu$G for Coma.  Studies of a large number of galaxy clusters with many RM probes per cluster are currently unfeasible given the sensitivity limits of available radio telescopes, combined with the small angular size of more distant clusters. This situation can be circumvented by obtaining RM probes through a sample of clusters, each having typically only one or two (bright) polarized radio sources. The consensus of these studies \\citep{ld82, kim91, gr93} is that cluster cores have a detectable component of RM, and many have field strengths at the microgauss level. Note that the Faraday study by \\citet{hoe89} does not find evidence for intracluster magnetic fields. This discrepancy is, however, likely due to the combination of small statistics and large impact parameters in their study.  For a few clusters with extensive cooling flows, high resolution Faraday rotation measure mapping of extended sources that are embedded within the cooling flow zones have, in combination with X-ray data, produced magnetic field strength estimates of 10--40 $\\mu$G, ordered on scales varying from 100 -- 0.5 kpc (see Taylor, Allen, \\& Fabian 1999, and references therein).  Lower limits to ICM magnetic field strengths in the range 0.1 -- 1 $\\mu$G level have been suggested from recent detections of both excess (over thermal) extreme ultraviolet (EUV) and hard X-ray (HEX) emission in some clusters. These field values would seem to {\\it prima facie} contradict the much higher values above {\\it if} the EUV excess emission is interpreted as inverse Compton (IC) scattering of $\\sim$ 100 MeV electrons \\citep{rug94}. The EUV detections \\citep{bowy1} in particular apply to only a few clusters having widespread extended synchrotron emission and/or they occur near to the cooling flow zone (as in M87). Spatial differentiation of high and low magnetic field regions in the ICM avoids the apparent contradiction by allowing the synchrotron and EUV IC emission to originate in low field regions, while high field regions, where the synchrotron energy loss time is short, would provide the major contribution to the Faraday rotation measures \\citep{ensslin1}. Further, EUV emission also appears not to be cluster wide, but concentrated to central sub-regions, and hence does not spatially correspond to the wider regions probed by the X-ray emission and Faraday RM measurements. The HEX excess in clusters can be plausibly understood as bremsstrahlung from a (probably shock heated) population of suprathermal electrons \\citep{ensslin1,dogiel1,sarazin00}. Thus, on balance, it appears that spatial differentiation of field regions and/or emission mechanisms other than IC scattering of CMB photons are more plausible where apparent contradictions in magnetic field estimates arise.  This paper concentrates on what we shall term ``normal'' Abell clusters {\\it i.e.\\ }those which have neither {\\it widespread} cooling flows nor strong synchrotron halos. The observations are aimed at estimating the strength and spatial extent of intracluster medium magnetic fields for a relatively homogeneous sample of 16 Abell clusters. Throughout this {\\it Letter}, we adopt $H_o=75\\,h_{75}\\ {\\rm km\\ s^{-1}\\ Mpc^{-1}}$, and $q_o$=0.5. ",
        "conclusions": "\\label{discussion}  These results confirm the widespread existence of magnetic fields in the central regions of non cooling flow clusters. The cluster-enhanced RM can generally be traced out to the periphery of the ROSAT-detectable ICM X-ray emission. The rotation measure distribution across our cluster sample drops from $\\sim$ 200 ${\\rm rad\\ m^{-2}}$ in the central regions (which may be an underestimate, see \\S 2) to the background level of $\\sim$ 15 ${\\rm rad\\ m^{-2}}$ at large radii.  Our new measurements of (1) the RM, (2) the intracluster electron density, (3) the magnetic field volume filling factor, and (4) the average tangling scale of the field enable us to estimate, even if only crudely, an important physical quantity $-$ the total energy in the ICM magnetic field for ``normal'', non-merging, relaxed clusters. For a 5$\\,h_{75}^{1/2}\\ \\mu $G magnetic field in the inner 500\\,$h_{75}^{-1}$ kpc sphere, the total magnetic energy is $E_{B} = 1.5 \\times 10^{61}({\\frac{r}{500\\ {\\rm kpc}}})^{3}(\\frac{B} {5\\ \\mu {\\rm G}})^{2}$\\,$h_{75}^{-2}\\ \\rm ergs$. The magnetic energy content of the ICM can then be compared to the total thermal energy content in the same cluster volume. The latter is (again taking constant values within a fiducial radius that is close to both the RM cutoff radius in Figure~\\ref{RRM} and the X-ray radius) $E_{th} = 6.4 \\times 10^{62}\\,n_{e}\\left ({\\frac{r}{500\\ {\\rm kpc}}}\\right )^{3} \\left (\\frac{T}{10^{8}\\ {\\rm K}}\\right ) \\,h_{75}^{-2}\\ {\\rm ergs}$, where $n_{e}$ is the intracluster electron density in units of $10^{-3}\\,h_{75}^{1/2}\\ {\\rm cm^{-3}}$. The ratio $ \\frac{E_{B}}{E_{th}}$ is of order 2.5\\%, and the possibility that $|B|$ may be underestimated due to limited radio resolution makes it possible that $ \\frac{E_{B}}{E_{th}}$ could be even higher. Even at the lower bound of 2.5\\% the ratio suggests that the magnetic energy provides a non-negligible fraction of the energy budget of the ICM.  We can now compare this approximate magnetic energy estimate with other sources of energy that are relevant for a cluster: the thermonuclear (stellar) energy released in all cluster-member galaxies, the gravitational binding energy released from AGN/accretion disks, the cluster gas binding energy, and the energy associated with past merger events. The available energy from stellar sources cannot be much greater than $\\sim 10^{62}$ ergs \\citep{voelk00}, {\\it i.e.} of comparable magnitude, and is thus insufficient to maintain the magnetic fields, barring a very high energy conversion efficiency. This means that thermonuclear energy can be ruled out as the primary source of intracluster magnetic field energy. It must therefore be ultimately derived from gravity. A single powerful AGN/accretion disk can be expected to inject approximately $10^{61}$ ergs over its lifetime into the ICM. Comparing the lifetime of the radio source with the cluster lifetime we expect $10^2$ sources to have injected a total of $\\sim 10^{62} - 10^{63}$ ergs into the ICM. This makes AGN/accretion disks an attractive possible source of the ICM field energy, as has been suggested by \\citet{colgate99}. Other sources of ICM magnetic energy are the gravitational binding energy of the cluster gas, which is of order $10^{64}$ ergs, and the energy associated with a past cluster merger event, which is of order $10^{63-64}$ ergs. This suggests that the magnetic energy possibly ``taps into'' a combination of the gravitational energy released in AGN/accretion disks over the lifetimes of clusters, and shearing and shocks associated with larger scale infall of matter as the clusters evolve.  Given that the cooling flows represent a late stage of cluster evolution and our ``global'' value of 5 $\\,h_{75}^{1/2} \\mu G$ for ICM zones within the clusters, but outside of cooling flow zones, we conclude that the cooling flow develops out of a medium whose field strength is already a significant fraction of what is seen in the cooling flow zones."
    },
    "0011/astro-ph0011554_arXiv.txt": {
        "abstract": "We present a BeppoSAX observation of the X--ray source \\axj that has been proposed as the neutron star associated  to the shell like-supernova remnant \\rxj. The image  at E$>$5 keV  shows the presence of a new source, \\sax, not resolved from  \\axj in previous observations. The improved positional accuracy obtained for \\axj confirms its possible identification with an early type star, also consistent with its soft thermal spectrum (kT$\\sim$1.3 keV). \\sax, with a  harder X-ray spectrum and  a higher X--ray to optical flux ratio, is a more likely candidate for the neutron star associated to the supernova remnant. ",
        "introduction": "The supernova remnant \\rxj  consists of a circular X--ray shell, with an angular diameter of $\\sim$2$^{\\circ}$, seen in projection against the much more extended soft X--ray emission of the Vela supernova remnant (Aschenbach 1998). \\rxj (also known as G266.2--1.2) has  attracted considerable interest since Iyudin et al. (1998) reported its possible detection in the $\\gamma$--ray  line   of $^{44}$Ti (1.156 MeV). This isotope,   produced in different types of supernovae, has a lifetime of $\\sim$90 years and should therefore be a good tracer for young supernova remnants (SNRs). The observed $\\gamma$--ray line flux and the relatively small angular dimensions of the remnant imply an age of only $\\sim$680 years and a small distance of $\\sim$200 pc (Aschenbach, Iyudin \\& Sch\\\"{o}nfelder 1999). Thus \\rxj could be the remnant of the closest supernova event to have occurred in recent historical times.  However, this interpretation poses   a few problems. For instance, the lack of records for an event that could have reached the  optical brightness of the full moon (for a Type Ia SN), and in any case was brighter than Venus (for an extremely sub-luminous SN), is puzzling. If a neutron star was formed, it should still be a bright source of thermal emission, but no  bright soft X--ray sources have been found within \\rxj in the ROSAT All Sky Survey data. The upper limits reported by Aschenbach et al. (1999) imply surface temperatures smaller than 3$\\times$10$^5$ K, assuming the low interstellar absorption expected for a  distance of only a few hundreds parsecs.  Recent   observations of \\rxj obtained with the ASCA satellite have been reported by Slane et al. (2001). These new data clearly indicate the non-thermal nature of the X--rays from the SNR shell.  The high absorption  suggests that \\rxj is at a distance of the order of $\\sim$1-2 kpc, well beyond the Vela SNR. Consequently an age of several thousands year is more likely. These conclusions are also supported by the fact that a new analysis of the COMPTEL data (Sch\\\"{o}nfelder et al. 2000) has reassessed the statistical significance of the $^{44}$Ti emission from the direction of \\rxj. Depending on the data selection and method of analysis, the excess is present at a level between $\\sim$2 and 4 $\\sigma$: if \\rxj is not emitting $^{44}$Ti $\\gamma$-rays there are no compelling reasons to require a particularly small distance and age.  The ASCA observations revealed also a central point source, surrounded by diffuse X--ray emission, that might be the  neutron star associated to \\rxj. Here we report on a recent observation pointed toward the central region   of \\rxj performed with the BeppoSAX X--ray satellite. ",
        "conclusions": "A recent BeppoSAX observation of  the central region of the supernova remnant \\rxj has led to the detection of two X--ray sources that can be considered good neutron star candidates. The softer source, \\axj was already observed with ROSAT and ASCA. Though unlikely (Slane et al. 2001), it cannot be excluded that the X--rays from this source are due to one (or both) of the two early type stars present in its error region. The newly discovered X--ray source, \\sax, does not have obvious stellar counterparts. It  has a harder spectrum and properties  consistent with a  neutron star with an age of a few 10$^4$ years at a distance of a few kpc.  Unless very unusual properties are invoked, both sources are difficult to explain as the neutron star associated to \\rxj, if the latter is really as young and nearby ($\\sim$680 years, d$\\sim$200 pc) as implied by the claimed detection of a $^{44}$Ti  $\\gamma$--ray line (Iyudin et al. 1998)."
    },
    "0011/astro-ph0011232_arXiv.txt": {
        "abstract": "Two {\\it ASCA} observations were made of two ultra-luminous compact X-ray sources (ULXs), Source~1 and Source~2, in the spiral galaxy IC~342. In the 1993 observation, Source~2 showed a 0.5--10 keV luminosity of $6 \\times 10^{39}$ ergs s$^{-1}$ (assuming a distance of 4.0 Mpc), and a hard power-law spectrum of photon index $\\sim 1.4$. As already reported, Source~1 was $\\sim 3$ times brighter on that occasion, and exhibited a soft spectrum represented by a multi-color disk model of inner-disk temperature $ \\sim 1.8$ keV. The second observation made in February 2000 revealed that Source~1 had made a transition into a hard spectral state, while Source~2 into a soft spectral state. The ULXs are therefore inferred to exhibit two distinct spectral states, and sometimes make transitions between them. These results significantly reinforce the scenario which describes ULXs as mass-accreting black holes. ",
        "introduction": "Arm regions of nearby spiral galaxies have long been known to harbour luminous point-like X-ray sources (Fabbiano 1989, 1998; Read et al. 1997; Roberts \\& Warwick 2000). Except for several identified with young supernova remnants, they remained unidentified in other frequencies. Being often time variable, these sources are thought to be accreting single objects; we hence call them ultra-luminous compact X-ray sources (ULXs; Makishima et al. 2000, hereafter Paper~I). With the X-ray luminosity reaching $10^{39-40}$ ergs s$^{-1}$, the ULXs are suspected to contain stellar black holes (BHs) of mass up to $\\sim 100~M_\\odot$ or more (e.g., Colbert et al. 1995; Fabbiano 1998; Immler et al. 1999; Roberts \\& Warwick 2000), so as not to violate the Eddington limit. However, this much exceeds the mass estimates ($5-15~ M_\\odot$; Tanaka \\& Lewin 1995) for accreting stellar BHs in the Galaxy and the Large Magellanic Cloud. There is no consensus on the presence of $\\sim 100~M_\\odot$ BHs, even though such medium-mass BHs have been proposed under slightly different context (Colbert \\& Mushotzky 1999; Matsumoto \\& Tsuru 1999; Ptak \\& Griffiths 1999). Furthermore, we do not find ULXs in the Milky Way or M31. Therefore, the BH scenario of ULXs remained inconclusive, and their nature has remained a big mystery.  A great leap has been achieved through spectroscopy of a dozen ULXs with {\\it ASCA} (Takano et al. 1994; Okada et al. 1998; Colbert \\& Mushotzky 1999; Mizuno et al. 1999; Mizuno 2000; Kotoku et al. 2000; Paper~I), as the spectra of majority of them have been reproduced by so-called multi-color disk model (MCD model; Mitsuda et al. 1984) that describes emission from an optically-thick standard accretion disk around a BH (Shakura \\& Sunyaev 1973). These results significantly reinforce the BH scenario of ULXs, since the MCD model can explain the soft X-ray spectra of Galactic/Magellanic black-hole binaries (BHBs) in ``soft (or high) state'' (e.g., Makishima et al. 1986; Ebisawa et al. 1993;  Dotani et al. 1997; Kubota et al. 1998; Feroci et l. 1999; Paper~I).  The scenario is nevertheless still clouded by a self-inconsistency, that the measured inner-disk temperature of ULXs, $T_{\\rm in} = 1.0-1.8$ keV, is too high for the implied high BH mass (Colbert \\& Mushotzky 1999; Paper~I). Although the scenario may be salvaged by invoking rapid BH rotation (Mizuno et al. 1999; Paper~I; Mizuno 2000; also Zhang et al. 1997b) and/or ``slim disk'' concept (Abramowicz et al. 1988; Watarai et al. 2000), it is important to assemble further observational clues.  Here, we report further evidence supporting the BH interpretation of ULXs, i.e., transitions of some ULXs between two distinct spectral states. These results are based on two observations with {\\it ASCA} (Tanaka et al. 1994) of two ULXs, both in the spiral galaxy IC~342 at a distance of 4.0 Mpc (Tully 1988).  \\placefigure{fig:images} ",
        "conclusions": "In the first {\\it ASCA} observation of IC~342, Source~1 exhibited a high luminosity exceeding $10^{40}$ ergs s$^{-1}$ (Table~1), prominent short-term variation, and a convex spectrum expressed with an MCD model of a high disk temperature, $T_{\\rm in} \\sim 1.8$ keV. In Paper~1, these features have been identified as prototypical ULX properties.  In the same observation, Source~2 showed a distinct hard spectrum expressed with a power-law of $\\Gamma \\sim 1.4$. Among the ULXs so far studied with {\\it ASCA}, another object, i.e., Source~A in NGC~1313 observed in 1993, exhibited a similarly hard spectrum with $\\Gamma \\sim 1.7$ (Petre et al. 1994; Colbert \\& Mushotzky 1999; Mizuno 2000); the remaining objects showed the MCD-type spectra (Paper~I). We hence presume that there are two subtypes of ULXs, a majority exhibiting the soft MCD-type spectra, and a minority showing the hard power-law spectra. However, from these results alone, we cannot tell whether or not these subtypes represent intrinsically distinct two source populations.  The second {\\it ASCA} observation made in 2000 revealed that Source~1 had made a soft-to-hard state transition, whereas Source~2 a hard-to-soft one, although their reciprocal behavior must be a chance coincidence. To visualize this, we present in Figure~\\ref{fig:deconvolved} their deconvolved spectra. These results immediately imply that; (1) the two ULXs must be accreting compact objects, to make such drastic spectral changes in 7 years; (2) the two ULX subtypes represent two well-defined states (soft state and hard state) of the same source population; (3) some (if not all) ULXs make transitions between the two spectral states; and (4) at least in the {\\it ASCA} band, the source is less luminous while in the hard state. These properties give a conclusive support to the BH interpretation of ULXs, because they are just what characterize the known BHBs (e.g., Maejima et al. 1984; Tanaka \\& Lewin 1995; Tanaka \\& Shibazaki 1996; Zhang et al. 1997a; Wilms et al. 2000). A similar (though less convincing) example may be the aforementioned Source~A in NGC~1313, because its spectrum turned very soft in 1995 ($\\Gamma \\sim 2.8$; Colbert \\& Mushotzky 1999; Mizuno 2000).  \\placefigure{fig:deconvolved}  Thus, the long-lasting puzzle of ULXs are being unveiled, but a series of new questions are arising in turn. Which is the right solution to the ``too high $T_{\\rm in}$''  problem? Are the inferred high BH masses of ULXs consistent with the current understanding of stellar BH formation? Why are ULXs absent in Our Galaxy and M31? Do they form a population distinct from the Galactic/Magellanic BHBs? In short, the ULX formation scenario is yet to be clarified.  The authors thank the {\\it ASCA} team members. The present work is supported in part by the Grant-in-Aid for Center-of-Excellence, No. 07CE2002, from Ministry of Education, Science, Sports and Culture of Japan."
    },
    "0011/astro-ph0011004_arXiv.txt": {
        "abstract": "We study the differential energy distribution of dark matter halos, carrying out cosmological $N$-body simulation. We give an analytical formula of the differential energy distribution of dark matter in the halos obtained by the numerical simulation. From the analytical formula we reconstruct the density profile described by the Navarro, Frenk, \\& White (NFW) profile. The NFW profile is consistent with the analytical formula of our fractional mass distribution. We find that a parameter in our analytical formula of differential energy distribution which is related with the slope of inner cusp of dark halo. We obtain the distribution function for the NFW profile which has sharp cut off at the high binding energy. We discuss physical reason of form of the analytical formula. ",
        "introduction": "It is interesting to understand how the density profiles of galaxies and clusters of galaxies have formed. Navarro, Frenk and White (1995, 1996, 1997; NFW) have shown in their N-body simulations of Cold Dark Matter (CDM) in the standard biased CDM and four power law spectra with indices $n = 0,-0.5,$ and $-1$, open CDM ($\\Omega_0 =0.1$) with power-law spectra ($n=0$ and $-1$), and $\\Lambda$CDM cosmology that density profiles of dark halos have an universal profile described as \\begin{eqnarray} \\rho(r) \\propto \\frac{1} {\\left(\\frac{r}{r_s}\\right) \\left( 1 + \\frac{r}{r_s} \\right)^2}. \\end{eqnarray} Several investigators have shown that the formula provides a good fit to their numerical results (Cole and Lacey 1996, Tormen, Bouchet \\& White 1997, Huss, Jain \\& Steinmmetz 1999, Thomas {\\it et al.} 1998). However, some other simulations (Fukushige \\& Makino 1997; Moore et al. 1998; Okamoto \\& Habe 1999, 2000) indicate that their density profiles have steeper inner cusp than the NFW profile. Jing (2000) gives his numerical results that steeper cusps of density profile are found in recent mergers and in dark halos with substructures, by a large set of high resolution cosmological simulations.  Subramanian, Cen and Ostriker (2000) discussed the general theoretical grounds of the density profiles. They discussed that there is a possible connection between slope of inner region of dark halo and the formation epoch and proposed that there is possible relation to  cosmological parameters.  However, physical reason of the NFW profile have not been explored, although it is very important to understand formation of galaxies and clusters of galaxies.  We study the differential energy distribution of the NFW profile, $d M/ d \\varepsilon$, where $\\varepsilon$ is the binding energy. The differential energy distribution was studied to understand virialization process of $N$- body system (van Albada 1982). Binney (1982) shows that $dM/d \\varepsilon$ of the elliptical galaxies is approximated by the Boltzmann factor, if surface brightness of elliptical galaxies obeys the de Vaucoureurs' $r^{1/4}$ law and their mass-to-light ratio is constant. It is interesting to study phase space structure of virialized objects with the NFW profile, since the differential energy distribution may give insight of the relaxation process (Binney \\& Tremain 1987).  In this paper we calculate formation of clusters of galaxies by $N$-body simulation and we study the differential energy distribution of the clusters in our numerical results. We find an analytical formula of the fractional mass distribution that is defined by the differential distribution divided by mass of the object, is fitted by an analytical formula. We show that the NFW profile is consistent with our analytical formula of the fractional mass distribution. Using the iteration method, we construct the density profile from the analytical formula to show how the slope of the cusp of the density changes with parameters in our analytical formula.  In the next section, we illustrate our numerical method. In the section 3, we present our numerical results. We show the iteration method and give results by this method in the section 4, and we summarize and discuss our results in the section 5. ",
        "conclusions": "We analyze the universal density profile of dark halo proposed by NFW and its differential energy distribution. Our main results are summarized as follows.  \\begin{enumerate} \\item We study the fractional mass function $N(\\varepsilon)$ for dark halo obtained by our numerical simulation and find its analytical formula which is the equation (\\ref{escort}). \\item We show that the NFW profile is given by the equation (\\ref{escort_m}). \\item We show that the slope of the cusp in the density profile changes with a value of the parameter $q$ in the analytical formula. \\end{enumerate}  We can regard that $N(\\varepsilon)$ shows the statistical property of the NFW profile. If the NFW profile is universal, $q = 0.6 - 0.7$ in equation (\\ref{escort}). Different $q$ makes slope of a cusp different. Since $q$ plays an important role, we should make clear what physical process determines $q$. Recent high resolution numerical simulation (Okamoto \\& Habe 1999, 2000) shows the steeper cusp, $\\rho \\propto r^{-1.5}$, than the the NFW profile. This profile corresponds to $q = 0.75 - 0.8$. Isothermal profile, $\\rho \\propto r^{-2}$, corresponds to $q = 1$.  We study $f(\\varepsilon)$ for the NFW profile. The formula of this is not isothermal one nor the King formula $f_K \\propto e^{\\varepsilon/\\sigma^2} -1$. $f(\\varepsilon)$ for the NFW profile have the energy cut off at the high end of $\\varepsilon$. We should study the reason why $f(\\varepsilon)$ has such a form. Lynden-Bell (1967) studied distribution function $f(\\varepsilon)$ of a virialized system. Maximizing the Boltzmann entropy of the system, resulting distribution is isothermal profile, $ \\rho \\propto r^{-2}$. In this case the system has infinite extend and infinite mass. This is not realistic for astronomical objects. Cosmological simulations have shown that galaxies and clusters of galaxies formed in these simulations have more rapid radial decline than isothermal in the outer part.  We note that the form of the distribution function relates with the cusp profile. For $q < 1$, there are the maximum binding energy $\\varepsilon_{{\\rm max}}$ and the sharp peak in $f(\\varepsilon)$ at $\\varepsilon_{{\\rm max}}$. Since there is $\\varepsilon_{{\\rm max}}$, phase space volume occupied by dark matter is limited than the isothermal profile. This may be the reason of the cusp profile different from $\\rho \\propto r^{-2}$. For $q < 1$, the distribution function $f(\\varepsilon)$ changes with $q$ values as shown in figure 8.  We note that the form of equations (\\ref{escort}) and (\\ref{escort_m}) are similar to the Tsallis' escort distribution, \\begin{eqnarray} \\frac{P(E,T^{\\prime})}{P(0,T^{\\prime})} = \\left[ 1 - ( 1- q) \\frac{E}{T^{\\prime}} \\right]^{q/(1-q)}, \\end{eqnarray} where $T^{\\prime}$ is temperature parameter and $q$ is entropic index (Tsallis, Mendes, \\& Plastino, 1998). Tsallis' non-extensive generalized statistics (Tsallis 1988) is paid attention in the area of statistics of a multi-fractal system. In non-extensive system (long-range microscopy memory, long range forces, fractral space time) the following generalized entropy has been proposed : \\begin{eqnarray} S_q = k \\frac{1 - \\sum_i p_i^q}{q-1} \\ \\ \\ \\left( \\sum_i = p_i = 1; \\ \\ \\ q \\in \\Re \\right), \\end{eqnarray} where $k$ is a positive constant. Optimization of $S_q$ yields, for the canonical ensemble, \\begin{eqnarray} &&p_i = Z_q^{-1} [1 - (1-q) \\varepsilon_i/T^{\\prime} ]^{1/(1-q)}, \\\\ &&Z_q \\equiv \\sum_i [1 - (1-q) \\varepsilon_i/T^{\\prime} ]^{1/(1-q)} \\end{eqnarray} and, when $q \\longrightarrow 1$, the Boltzmann-Gibbs result is recovered. In this statistics an expected value of any physical variable is given by the Tsallis' escort distribution : \\begin{eqnarray} \\left< A \\right>_q = \\frac{ \\sum_i p_i^q A_i}{\\sum_j p_j^q}, \\end{eqnarray} where $\\{A_i\\}$ are the eigenvalues of an arbitrary observable $A$. The escort distribution, $P_i = p_i^q/\\sum_j p_j^q$, is similar to the form of equation (\\ref{escort}) and (\\ref{escort_m}).  It is expected that long range interaction makes the system non-extensive. Our case may be the one of this cases. Our results suggests that differential energy distribution of collisionless particles is described by the escort distribution of the Tsallis statistics. This may indicate that energetic process of collisionless particles must be stochastic process.  In the Tsallis' escort distribution function, there is the maximum value of $\\varepsilon$ for $p_i > 0$ for $0 < q < 1$. This case is called superextensive. The shallower cusp profile in the the NFW profile shows the superextensive property of the dark matter distribution.  We show that the NFW profile corresponds to $q = 0.6-0.7$. Lavagno {\\it et al.} (1998) have recently shown that fraction of peculiar velocity of cluster of galaxies (Bahcall \\& Oh 1996) is well explained by the Tsallis escort integral, \\begin{eqnarray} P(> v) = \\frac{\\int_v^{v_{max}} (1 - (1-q)(v/v_0)^2)^{q/(1-q)}dv} {\\int_0^{v_{max}} (1 - (1-q)(v/v_0)^2)^{q/(1-q)}dv}. \\end{eqnarray} They obtained $q=0.23$ to fit the fraction of peculiar velocity of cluster of galaxies and is smaller than in our case. There is a conjecture that $q$ of system approaches unity when the system proceeds relaxation (Tsallis 1999). Our results in which $q = 0.6-0.7$ are consistent with this conjecture, since dark matter distribution in a cluster of galaxies is more relaxed system than the large scale motion of clusters of galaxies. The value of $q$ must be related with the degree of relaxation of collisionless particles. $q = 1$ corresponds to isothermal distribution in the Tsallis statistics. The gravitational system like clusters of galaxies with the universal profile may be non-extensive because they formed recently.   We should study the reason why dark halo has the value of $q = 0.6-0.7$ in the hierarchical clustering scenarios. It is interesting to study the differential energy distribution of self-gravitational system formed in a circumstance without hierarchical clustering to make clear mechanism what determines $q$ of the gravitational system."
    },
    "0011/astro-ph0011374_arXiv.txt": {
        "abstract": "Despite previous intensive ground--based imaging and spectroscopic campaigns and wide-band HST imaging of the $z=0.927$ QSO 3C336 field, the galaxy that hosts the damped \\Ly\\ system along this line--of--sight has eluded detection. We present a deep narrow-band \\Ha\\ image of the field of this $z_{abs}=0.656$ damped \\Ly\\ absorber, obtained through the F108N filter of NICMOS 1 onboard the {\\it Hubble Space Telescope}. The goal of this project was to detect {\\it any} \\Ha\\ emission 10 times closer than previous studies to unveil the damped absorber. We do not detect \\Ha\\ emission between 0.05\\arcsec\\ and 6\\arcsec\\ (0.24 and 30 $h^{-1}$~kpc) from the QSO, with a 3$\\sigma$ flux limit of $3.70\\times 10^{-17} h^{-2}$ \\flux\\ for an unresolved source, corresponding to a star formation rate (SFR) of $0.3 h^{-2}$~\\mpy. This leads to a 3$\\sigma$ upper limit of 0.15 \\mpy~kpc$^{-2}$ on the SFR density, or a maximum SFR of 1.87~\\mpy\\ assuming a disk of $4$~kpc in diameter. This result adds to the number of low redshift damped \\Ly\\ absorbers that are {\\it not} associated with the central regions of Milky-Way-like disks.  Damped \\Ly\\ absorption can arise from high density concentrations in a variety of galactic environments including some that, despite their high local \\HI\\ densities, are not conducive to widespread star formation. ",
        "introduction": "QSO absorption lines provide a powerful  approach to studying the history of galaxies.  The nature of damped \\Ly\\ absorbers (DLAs) towards background QSOs has been an ongoing debate for more than a decade.  \\citet{wolfe86,wolfe95} proposed that DLAs are large progenitors of today's massive spiral disks.  Evidence in support of this interpretation includes the measurement of absorption line velocity profiles that are consistent with those expected from lines of sight intercepting rotating, thick gaseous disks \\citep{prochaska}. However, recent theoretical simulations of galaxy formation showed that a large range in structures (e.g. halo gas clouds) and morphologies, rather than a single uniform type of galaxy, can give rise to DLAs (e.g. \\citep{katz,haehnelt,mcdonald}).  In this case, they could be low surface brightness (LSB), gas rich, dwarf galaxies as proposed by \\citet{tyson} and more recently by \\citet{jimenez}.   DLAs make up the largest reservoir of neutral hydrogen (\\HI) at high redshift \\citep{wolfe86,lanzetta,rao00}.  There are hundreds of DLAs, defined as absorbers with atomic hydrogen column densities $N(\\HI)$ greater than $ 2 \\times 10^{20}$~{\\cmsq}, known up to redshifts $z\\sim4$.  From the redshift distribution of the measured column densities of the damped systems, the evolution of neutral gas density ($\\Omega_{DLA}(z)$) can be measured \\citep{lanzetta,lanzetta95,wolfe95,storrie96,rao00}.  The analysis of the variation of $\\Omega_{DLA}(z)$ with redshift can provide another measurement of the evolution of the star formation rate in the Universe \\citep{pei}.  \\citet{rao00} find little evolution in $\\Omega_{DLA}(z)$ from $z=4$ to $z=1$.  Over this interval we might have expected a decrease in $\\Omega_{DLA}$ as the gas is converted into stars in order to maintain the observed constant star formation rate \\citep{steidel99}.  This led \\citet{rao00} to conclude that DLAs and the galaxies that dominate the star formation density are different populations.   Furthermore, it is difficult to reconcile the low metallicities of high-redshift DLAs (typically 1/10 of solar at $z\\sim2.5$, \\citet{pettini}) with the higher metallicities of stars in galaxies today: no chemical evolution is seen in DLAs from $z=3.5$ to $z=0.3$ \\citep{pettini99}, which indicates that DLAs do not necessarily trace the population responsible for the bulk of star formation.  Finally, high resolution Keck spectra of three QSOs by \\citet{pettini00} indicate that DLAs ($z<1.0$) have heterogeneous chemical properties.  Another approach to understanding DLAs is to compare them with current local \\HI\\ surveys \\citep{briggs,zwaan,rosenberg}, which sample neutral gas clouds and galaxies perhaps analogous to distant DLAs (of course, high redshift and low redshift DLAs may be produced by very different galaxy populations).  Current results \\citep{zwaan,rao00} indicate that local \\HI\\ samples contain a much smaller fraction of high column densities ($N_{\\HI} >10^{21}$~cm~$^{-2}$) than both low and high redshift DLAs.  These results imply strong evolution at the highest column densities.  However, estimates of the local column density distribution rely on many caveats, such as the true nature of the debated local \\HI\\ mass function.  Beam-smearing in these \\HI\\ surveys may also lead to underestimates of the local number of high column density systems, particularly if such regions are physically small.   Prior to 1993, most observations of DLAs were restricted to high redshift systems ($z>1.8$) since the \\Ly\\ line is in the rest--frame UV.  Previous attempts to detect emission from DLAs have concentrated on \\Ly\\ (e.g., \\citet{lowenthal95}; see \\citet{roche} for a summary), which is expected to be a signature of a star-forming region, although its emission may be suppressed by dust extinction.  Of at least ten DLAs at $z\\ga 2$ searched for \\Ly\\ emission, only a very few show confirmed detections \\citep{djorgovski,fynbo00}.  Ground based surveys (photometric and spectroscopic) for {\\Ha} emission around $z>2$ DLAs have also been mostly unsuccessful \\citep{bunker,teplitz}, except in some cases \\citep{bechtold,mannucci}.  Only recently, with the advent of rest-UV spectroscopy from {\\it Hubble Space Telescope (HST)}, have data on intermediate redshift $0.3<z<1.8$ DLA systems become available \\citep{lebrun,boisse}.  These DLAs display a wide range of morphologies and surface brightnesses.  Using {\\it NICMOS-NIC2}, \\citet{kulkarni} reported a possible {\\Ha} detection of a $z_{abs}=1.89$ DLA towards LBQS 1210+1731 and \\citet{pettini00} detected an edge-on low luminosity $L\\sim 1/6 L^*$ galaxy 10~kpc away from the QSO 0058+019 ($z_{abs}=0.612$) using {\\it WFPC2}.  In addition, ground based observations such as WIYN images of three low redshift DLAs ($z < 0.3$) show dwarf and/or low surface brightness hosts, with confirmed redshifts \\citep{rao98,lane98}.   In contrast, the larger class of {\\MgII} absorbers with $3\\times 10^{17} \\leq N({\\HI})\\leq 2\\times 10^{20}$~{\\cmsq} (Lyman limit systems) are almost always associated with fairly luminous ($L_K>0.05L_K^*$) galaxies, i.e. within $35 h^{-1}$~kpc \\citep{steidel95}.  It was once thought that these different classes of absorbers sample different cross sectional regions of broadly the same galaxies \\citep{steidel93}, with the DLAs associated with the inner, denser regions.  This is only partially true.  Dwarf and LSB galaxies apparently can produce DLA absorption at low impact parameters; however, they do not contribute a significant cross--section for Lyman limit absorption.  These issues motivate the current attempts to image DLAs at low impact parameter and, at the same time, might explain why most previous attempts have failed to reveal gas-rich spirals.  In this paper, we present an {\\it HST} study of an intermediate redshift DLA ($z_{abs}=0.656$) along the line-of-sight towards QSO 3C336 ($z$=0.927) that has, so far, eluded detection entirely despite extensive ground-based searches \\citep{steidel}.  This QSO line-of-sight is one of the richest known for $z<1$ absorption line studies, with $6$ metal line systems in the interval $0.317 < z < 0.892$.  For that reason this quasar field was the target of both a very deep $24000$~s {\\it HST--WFPC2} image and a $2160$~s Keck/NIRC image \\citep{steidel}.  Five galaxies associated with the metal line systems were identified in the WFPC2 image and their redshifts confirmed spectroscopically using the Keck/LRIS \\citep{steidel}.  The only unidentified absorber is a DLA at $z=0.656$ with $N(\\HI)=2\\times 10^{20}$~{\\cmsq} and $[Fe/H]=-1.2$.  There is no galaxy detected with $L>0.05L_K^*$ near the QSO line-of-sight and as close as 0.5\\arcsec ($\\sim 2 h^{-1}$~kpc).  Two unlikely candidate galaxies exist at large impact parameters from the QSO. The first is a relatively faint, $M_K =-21.43$ ($m_K=20.77$), late type spiral located 14.3{\\arcsec} ($\\sim 65 h^{-1}$~kpc) NE of the QSO.  Taking into account the estimated disk inclination and the position angle, \\citet{steidel} estimated it would require a disk extent of at least $\\sim 120 h^{-1}$~kpc to intercept the QSO line--of--sight.  The second candidate is a galaxy without a confirmed redshift.  If one assumes this galaxy is at the redshift of the DLA, it would have an impact parameter of $41 h^{-1}$~kpc and $L_B=0.04L_B^*$ (similar to that of the SMC).  This raises the question of whether DLA absorption can arise in dense \\HI\\ regions far from the centers of galaxies, perhaps in regions that have little or no current star formation, as seen locally in mergers \\citep{hibb99} and as pointed out by \\citet{rao00}.  Alternatively, a separate absorbing galaxy could be situated beneath the QSO on the plane of the sky.   The goal of this project was to detect {\\it any} \\Ha\\ emission as close as $\\sim 0.05$\\arcsec ($0.24~h^{-1}$~kpc) of the QSO to test further the alternate hypothesis. \\Ha\\ at the redshift of this DLA matches one narrow--band filter of the {\\it HST-NICMOS} camera~1 and, therefore,  enables us to put strong constraints on the SFR of the absorber. We can already say there is no $L^*$ spiral galaxy close to the line-of-sight, for \\citet{steidel} did not find anything brighter than $0.05 L_K^*$. On the other hand, one might expect a dwarf or LSB galaxy with significant star formation such as found by \\citet{lebrun} and \\citet{rao00} for other DLAs.  In the next section, we describe the observations and the reduction of the data.  The results are given in section \\ref{sec_results}, and we compare them with previous studies in section \\ref{discussion}. Our conclusions are presented in section \\ref{conclusions}. Throughout this paper, we adopt $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=100 h$ km s$^{-1}$~Mpc$^{-1}$; thus 1\\arcsec\\ at $z=0.656$ corresponds to $4.85 h^{-1}$~kpc. ",
        "conclusions": "\\label{conclusions}  We can summarize previous observational results on DLAs as follows: (i) DLAs show little evolution of $\\Omega(z)$ from $z=4$ to $z=1$ \\citep{rao00}; (ii) DLAs show little metallicity evolution \\citep{pettini99} from $z=3$ to $z=0.3$; (iii) $z<1$ DLAs have heterogenous chemical properties \\citep{pettini00}; (iv) DLAs have various morphologies \\citep{lebrun,kulkarni,rao00} (although it is possible that at least {\\it some} DLAs are large disks like the galaxies found by \\citet{lebrun} and \\citet{pettini00}).  These results are inconsistent with the standard, DLA/{\\HI} disk paradigm\\citep{wolfe95} for they indicate that DLAs (1) do not participate in the overall chemical enrichment of the universe and hence do not trace star formation; and (2) are not characteristic of a particular type of galaxy.  Rather they are merely characteristic of a particular type of region: namely one with a large neutral column density (e.g. \\citet{kher96}).  It is quite reasonable, following the results of  \\citet{rao00}, to assume that this applies to higher redshift DLAs as well.  Nearly all space-based observations of low redshift DLAs have revealed star formation in some sort of galaxy.  The $z=0.656$ system toward 3C336 is an exception.  Our lack of detection of any source with SFR greater than 0.28$h^{-2}$~\\mpy, or $0.15$~\\mpy~kpc$^{-2}$, demonstrates that DLAs with very little star formation can exist.   As evidence mounts against the standard, DLA/{\\HI} disk scenario, we must address the question: ``What is a DLA?''. Possibilities not yet ruled out include knots of {\\HI} hundreds of kpc away from the main galaxy as seen locally \\citep{hibb99}; small compact dwarfs or LSBs; or photo-dissociation regions that may produce much of the {\\HI} currently observed in galaxy disks \\citep{smith}. In some cases, both a cold (hundreds of K) and a warm (thousands of K) neutral medium are found along the line of sight through a DLA, based upon analysis of 21--cm emission profiles \\citep{lane}, while in others the warm phase dominates.  Our limits on \\Ha\\ emission from the DLA towards 3C336 set the tightest constraints yet on star formation in a compact absorber at intermediate redshift, and probe closer to the QSO line-of-sight than any previous study.  Our non-detection adds to the mounting evidence that low redshift DLAs are made of galaxies of diverse morphologies, luminosities and surface brightnesses rather than a uniform population of luminous disk galaxies."
    },
    "0011/astro-ph0011142_arXiv.txt": {
        "abstract": "Gravitational lenses offer the possibility of accurately determining the Hubble parameter ($H_0$) over cosmological distances, and B0218+357 is one of the most promising systems for an application of this technique. In particular this system has an accurately measured time delay ($10.5\\pm0.4$~d; Biggs et al. 1999) and preliminary mass modelling has given a value for $H_0$ of 69$^{+13}_{-19}$\\,km\\,s$^{-1}\\,$Mpc$^{-1}$. The error on this estimate is now dominated by the uncertainty in the mass modelling. As this system contains an Einstein ring it should be possible to constrain the model better by imaging the ring at high resolution. To achieve this we have combined data from MERLIN and the VLA at a frequency of 5~GHz. In particular MERLIN has been used in multi-frequency mode in order to improve substantially the aperture coverage of the combined data set. The resulting map is the best that has been made of the ring and contains many new and interesting features. Efforts are currently underway to exploit the new data for lensing constraints using the LensClean algorithm \\cite{kochanek92}. ",
        "introduction": "The gravitational lens system B0218+357 \\cite{patnaik93} has long been recognised as an excellent system for measuring the Hubble constant ($H_0$) over cosmological distances using the method of Refsdal (1964). The system has a simple morphology comprising two compact images (A and B) of a flat-spectrum radio core and a radio Einstein ring (Figs~\\ref{mer1}, \\ref{mer3} and~\\ref{mervla}). There is also a kpc-scale radio jet that extends south from the bottom of the Einstein ring (seen most clearly in Fig~\\ref{vla}). Both lens and source redshifts are well determined, at 0.6847 \\cite{browne93} and 0.96 \\cite{lawrence96} respectively. In addition, the time delay between the two compact components has been accurately measured ($10.5\\pm0.4$~d) from VLA monitoring \\cite{biggs99}. Preliminary modelling of the mass distribution in the lensing galaxy using a Singular Isothermal Ellipsoidal (SIE) parameterization \\cite{kormann94} gave a value for $H_0$ of 69$^{+13}_{-19}$\\,km\\,s$^{-1}\\,$Mpc$^{-1}$. The quoted error is a 95~per~cent confidence limit (statistical). Constraints on this model come from the VLBI substructure of the compact images (Patnaik, Porcas \\& Browne 1995) as well as the flux density ratio measured from the VLA monitoring. A free parameter in the model is the position of the lensing galaxy centre which is poorly determined from {\\em Hubble Space Telescope (HST)} optical and infra-red observations. Leh\\'{a}r et al. (2000) have pointed out that the uncertainty in the position of the lensing galaxy implies an uncertainty in $H_0$ that is much greater than that given in Biggs et al. (1999).  A potential source of model constraints in B0218+357 is the radio brightness distribution in the Einstein ring which has been neglected until now. Einstein rings have proved to be a valuable source of constraints in other lens systems (e.g. MG~1131+0456 and PKS~1830-211) as they probe the mass distribution in the lensing galaxy at many points, thus providing many more constraints than are available from lenses that do not contain images of large-scale extended structure \\cite{kochanek90}. Techniques that have been developed to optimise lens models using radio maps of Einstein rings include the Ring Cycle algorithm \\cite{kochanek89} and LensClean \\cite{kochanek92}.  The Einstein ring in B0218+357 is believed to be an image of part of the extended emission of the kpc-scale radio jet. It is also the smallest known Einstein ring with a diameter of only 335~milliarcsec (mas), equal to the separation between the two compact components. Because of its small size, producing reliable maps of the brightness distribution in the ring that have high resolution and high sensitivity is difficult. The previous best map was made at a frequency of 5~GHz with the MERLIN array \\cite{patnaik93}, but due to the incomplete aperture coverage of the observations no attempt was made to interpret details of the Einstein ring in a lensing context. This paper presents a new map of B0218+357 made from data obtained with both MERLIN (in multi-frequency mode) and the VLA at frequencies around 5~GHz. The excellent aperture coverage when these data are combined has allowed us to produce the best map of this system at this frequency to date. Work is currently underway to exploit the new map for model constraints (Wucknitz et al., in preparation). ",
        "conclusions": "In this paper we have presented a new image of the gravitational lens system B0218+357 at a mean frequency close to 5~GHz. This has combined data from both the VLA and MERLIN, the latter in multi-frequency mode observing at three separate frequencies. The resulting data set has excellent $(u,v)$ coverage and sensitivity, as well as resolution high enough ($\\sim50$~mas) to easily resolve the Einstein ring. The new map is the best made of this system.  The motivation for making this map was to improve the lens model for this system and hence to improve the estimate of $H_0$. This currently stands at 69$^{+13}_{-19}$\\,km\\,s$^{-1}\\,$Mpc$^{-1}$ \\cite{biggs99}, but there are only limited constraints on the mass model available from the VLBI substructure of the compact images and the centre of the lensing galaxy is poorly determined. The next stage will be to exploit the new image for lensing constraints. The results of applying the Lens Clean algorithm will allow us to improve substantially the estimate of $H_0$ from the time delay."
    },
    "0011/astro-ph0011468_arXiv.txt": {
        "abstract": "BOOMERanG has recently resolved structures on the last scattering surface at redshift $\\sim$ 1100 with high signal to noise ratio. We review the technical advances which made this possible, and we focus on the current results for maps and power spectra, with special attention to the determination of the total mass-energy density in the Universe and of other cosmological parameters. ",
        "introduction": "A wide class of cosmological models predicts a harmonic series of peaks in the power spectrum of the Cosmic Microwave Background (CMB). These are the imprint of the acoustic oscillations inside the horizon in the primeval plasma. At recombination ($z \\sim 1100$, $t \\sim 300000 h^{-1}$ years) the acoustic horizon subtends an angle of roughly $1^o$, corresponding to a first peak at multipoles $\\ell \\sim 200$ in the spherical harmonic expansion of the CMB temperature. A high confidence measurement of this peak came from analysis of the high quality image of about 3$\\%$ of the sky obtained in the long duration flight of BOOMERanG (de Bernardis et al. 2000; Lange et al. 2000). An important confirmation arrived soon after from the MAXIMA experiment (Hanany et al. 2000; Balbi et al. 2000). Here we focus on the BOOMERanG experiment, which has produced a wide ($\\sim 1800$ square degrees), faithful image of the CMB at angular scales smaller than 10$^o$. ",
        "conclusions": "The data from BOOMERanG strongly support two of the main predictions of the simplest inflationary models: a nearly flat geometry of the Universe and a nearly scale-invariant spectrum of primordial density fluctuations. The detected density of baryons is from higher to consistent with that predicted from BBN, depending on the assumptions for other parameters. Combined with observations of different nature the BOOMERanG data produce significant detections of $\\Omega_\\Lambda$ and $\\Omega_m$. These conclusions are all confirmed by a joint analysis of the COBE, BOOMERanG and MAXIMA-1 data (Jaffe et al. 2000). Further insights are expected from the analysis of the data from the other detectors in BOOMERanG, from the data of the second flight of MAXIMA, and from the data of several interferometric CMB experiments underway."
    },
    "0011/gr-qc0011105_arXiv.txt": {
        "abstract": "We examine the behavior of an anisotropic brane-world in the presence of inflationary scalar fields. We show that, contrary to naive expectations, a large anisotropy does not adversely affect inflation. On the contrary, a large initial anisotropy introduces more damping into the scalar field equation of motion, resulting in greater inflation. The rapid decay of anisotropy in the brane-world significantly increases the class of initial conditions from which the observed universe could have originated. This generalizes a similar result in general relativity. A unique feature of Bianchi I brane-world cosmology appears to be that for scalar fields with a large kinetic term the initial expansion of the universe is {\\em quasi-isotropic}. The universe grows more anisotropic during an intermediate transient regime until anisotropy finally disappears during inflationary expansion. ",
        "introduction": "Observations of galaxies, QSO's and the cosmic microwave background appear to indicate that we live in a universe which is remarkably uniform on very large scales. Yet the homogeneity and isotropy of the universe is difficult to explain within the standard relativistic framework since, in the presence of matter, the class of solutions to the Einstein equations which evolve towards a FRW universe is essentially a set of measure zero~\\cite{ch73}. The above statement is however only true for space-times containing `normal' matter satisfying `energy conditions' which ensure that (i) negative pressures can never grow so large as to dominate the energy density: $T_{00} \\geq \\vert T_{ij}\\vert$; (ii) the sum of the principle pressures of the fluid must be non-negative: $\\sum_{i=1}^3 T_{ii} \\geq 0$. The inflationary scenario, based as it is on a form of matter which violates these energy conditions, radically alters the above picture. Indeed, as demonstrated in~\\cite{wald,star83,ms86,tw86,js86,rellis86}, a large class of spacetimes both homogenize and isotropize under the influence of an effective cosmological $\\Lambda$-term. Thus the inflationary scenario can successfully generate a homogeneous and isotropic FRW-like universe from initial conditions which, in the absence of $\\Lambda$, would have resulted in a universe far removed from the one we live in today.  Recently there has been a great deal of interest in a cosmological scenario in which matter fields are confined to a 3 dimensional `brane-world' embedded in a higher dimensional `bulk' space~\\cite{rs}. This higher-dimensional cosmology generalizes the standard Kaluza-Klein picture by allowing the presence of large or even infinite non-compact extra dimensions. The issue of inflation on the brane was investigated in~\\cite{mwbh}, where it was shown that on an FRW brane in 5-dimensional anti de Sitter space, extra-dimensional effects are conducive to the advent of inflation (see also~\\cite{cll}). In this paper we address the kindred issue of anisotropic initial conditions. We demonstrate that even very large initial anisotropy cannot prevent brane-world inflation from occurring, thus generalizing a previous result in general relativity~\\cite{ms86,tw86,js86,rellis86}. On the contrary, for a large class of initial conditions, the presence of anisotropy actually {\\em increases} the amount of inflation. Thus a scalar field dominated universe can eventually isotropise and inflate, even if its expansion was very anisotropic to begin with. A unique feature of brane cosmology is that the effective equation of state at high densities can become ultra stiff. Consequently matter can overwhelm shear for equations of state which are stiffer than dust, leading to quasi-isotropic early expansion of the universe in such cases. ",
        "conclusions": "Our results illustrated in Figs.~1 and 3 show the kinetic, potential and `anisotropy' energy densities plotted as functions of time. The associated dimensionless shear parameter $\\Omega_{\\rm shear} = \\sigma^2/6H^2$ is shown in Figs.~2 and 4 and the expansion factor in Fig. 5.   \\begin{figure}[tbh] \\includegraphics[width=3.3in]{Engy1.ps} \\caption{Same as Fig.~1, but with a larger initial value of the field velocity, $\\dot{\\phi} = -10^{-2}$ and $\\rho/\\lambda = 5 \\times 10^5$. } \\end{figure}   We find that the influence of anisotropy on the kinetic energy is particularly strong if $\\Sigma^2/a^6\\kappa^2 \\gg {\\dot\\phi}^2$ initially. In this case the kinetic term drops to a very small value, then rises after the anisotropy has disappeared, gradually approaching its asymptotic `slow-roll' value \\begin{equation}\\label{slow} {\\dot \\phi} \\simeq -V'/3H ~\\Rightarrow ~ {\\dot\\phi}^2 \\simeq \\frac{\\lambda}{3\\pi}\\left(\\frac{M_{\\rm p}}{\\phi}\\right)^2\\,. \\end{equation} We assume here that the potential has the simple `chaotic' form $V ={1\\over2} m^2\\phi^2$, for which the standard inflationary slow-roll condition is ${\\dot\\phi}^2 \\simeq m^2 M_{\\rm p}^2/16\\pi^2$. Comparing with Eq.~(\\ref{slow}), we find that dependence on $m^2$ has been replaced in the brane scenario by $\\lambda/\\phi^2$. Thus the kinetic energy does not remain constant but gradually increases as the field amplitude decreases during slow-roll. We find that the decay of anisotropy is generically accompanied by a corresponding decrease in the kinetic energy of the scalar field. This effect leads to greater inflation; see Fig.~5.  \\begin{figure}[h] \\includegraphics[width=3.3in]{anis2.ps} \\caption{The evolution of the dimensionless shear parameter $\\Omega_{\\rm shear} = \\sigma^2/6H^2$ is shown as a function of time for the model in Fig. ~3. We find that the early and late-time expansion of the universe is isotropic, but the shear term dominates during an intermediate anisotropic stage.} \\end{figure}  Examining Eq.~(\\ref{f}) closely we find that the effective equation of state of matter when $\\rho \\gg \\lambda$ is $w_{\\rm eff} = 2w + 1$. Consequently for matter with $w = p/\\rho > 0$ the approach to the initial singularity is matter dominated and not shear dominated, due to the predominance of the matter term $\\rho^2/2\\lambda^2$ relative to the shear term $\\Sigma^2/a^6$. (Within the framework of scalar field models this corresponds to ${\\dot\\phi}^2 > 2V(\\phi)$.) The fact that the density effectively grows faster than $1/a^6$ for $w > 0$ is a uniquely brane effect. Within the framework of general relativity such behaviour is clearly impossible since it would demand an ultra-stiff equation of state $w > 1$. The expansion of the early universe is therefore characterised by three successive expansion epochs during which the Bianchi I model experiences: (i) initial quasi-isotropic expansion $\\sigma^2/H^2 \\to 0, t \\to 0$; (ii) transient anisotropy dominated regime $\\sigma^2/H^2 \\sim 1 $; (iii) anisotropy dissipation $\\sigma^2/H^2 \\to 0, t \\to \\infty$. These three stages are illustrated in Fig ~4. For smaller values of ${\\dot\\phi}^2$ (corresponding to $w < 0$) stage (i) is absent and the shear decreases monotonically from a large initial value; see Fig. ~2.  \\begin{figure}[tbh] \\includegraphics[width=3.3in]{inflate.ps} \\caption{The evolution of the scale factor with time is shown for the model in Fig.~3. The curves shown in increasing amplitude from bottom to top correspond to: (a) $\\Sigma=0$, $1/2\\lambda = 0$ (the general relativity case); (b) $\\Sigma=0$, $1/2\\lambda = 10^{10}$ and (c) $\\Sigma=1$, $1/2\\lambda = 10^{10}$. We see that the effect of both the extra-dimensional term and cosmic shear is to produce more inflation.} \\end{figure}   The decay of shear and the accompanying isotropization of the universe significantly increases the class of initial conditions from which the present universe could have originated.    \\[ \\] {\\bf Acknowledgements:}\\\\  We would like to thank Alexei Starobinsky and Alexey Toporenskij for helpful comments which led to improvments in the manuscript. Discussions with Naresh Dadhich are also acknowledged.\\\\"
    },
    "0011/astro-ph0011433_arXiv.txt": {
        "abstract": "We analyze the temporal behavior of Network Bright Points (NBPs) searching for low-atmosphere signatures of flares occuring on the magnetic network. We make use of a set of data acquired during coordinated observations between ground based observatories (NSO/Sacramento Peak) and the MDI instrument onboard SOHO. Light curves in chromospheric spectral lines show only small amplitude temporal variations, without any sudden intensity enhancement that could suggest the presence of a transient phenomenon such as a (micro) flare. Only one NBP shows spikes of downward velocity, of the order of 2 -- 4 Km/s, considered as signals of compression associated to a (micro) flare occurrence. For this same NBP, we also find a peculiar relationship between the magnetic and velocity fields fluctuations, as measured by MDI. Only for this point the B - V fluctuations are well correlated, suggesting the presence of magneto-acoustic waves propagating along the magnetic structure. This correlation is lost during the compression episodes and resumes afterward. An A6 GOES soft X-ray burst is temporally associated with the downward velocity episodes, suggesting that this NBP is the footpoint of the flaring loop. This event has  a total thermal energy content of about $10^{28}$ erg, and hence belongs to the microflare class. ",
        "introduction": "High resolution observations taken by, e.g., YOHKOH and SOHO during the solar minimum years have revealed small scale transient phenomena occurring everywhere on the solar surface. Their presence could have important consequences on the issue of coronal heating: depending on their energy content and distribution, and rate of occurrence, they could provide enough energy to heat the corona (see reviews in \\opencite{ulm91}). Moreover, if these transient events represent the low energy tail of the broader solar flare family, their study could more easily identify the ``basic'' flare mechanisms, not biased by global collective effects as in the case of major solar flares.  Impulsive variations of the emission in the corona or transition region, that could be defined as low-energy flares, have long been observed in both ``quiet Sun'' regions, within magnetic network structures (\\opencite{gol74}, \\opencite{por87}), and active regions \\cite{shim95}. % In this paper we will focus on the characteristics of the brightenings associated to the magnetic network.  \\inlinecite{kru97} observed some small scale, transient soft X-ray brightening events associated to network magnetic bipoles. These events are spatially and temporally associated with microwave emission, and show several similarities with more energetic flares, such as the same ratio between the total energy emitted in soft-X and radio frequencies, etc.\\ However, the total observed radiative losses for these events are more than one order of magnitude smaller than previously reported X-ray events. For these brightenings they introduced the term {\\it soft X-ray network flares}.  Recent SOHO results, based on MDI, EIT, SUMER and CDS observations, have further demonstrated the highly dynamic nature of the ``quiet magnetic network''. Many observations confirm the existence of EUV brightening events in small-size structures. \\inlinecite{harr97} and \\inlinecite{harr99} observed rapid variations in the emission of transition region lines, associated with junctions of the network magnetic field. The emission characteristics of these brightenings, called blinkers, suggest that they are a class of phenomena distinct from the network flares. Spectrographic observations have shown bursts of highly Doppler-shifted emission, called transition region explosive events (among others, \\opencite{der94}; \\opencite{inn97}; \\opencite{cha98}). \\inlinecite{cha00} explored the relationships between blinkers and explosive events. They concluded that although the two phenomena appear different because of the differences in size and duration, they seem to be closely related to each other, for explosive events occur at the edges of blinkers and at the same time. Magnetic reconnection in different magnetic geometries is invoked to explain the observed differences.  \\inlinecite{ben99} and \\inlinecite{kru00} searched for EUV  and radio emission variations in non active regions associated to the magnetic network. They observed faint coronal and transition region fluctuations sharing some general properties with more energetic flares (such as the delay between emission originating at different atmospheric layers) and interpreted them as  micro- or nano-flares occurring onto the magnetic network.  The investigations mentioned above refer predominantly to coronal and transition region spectral signatures and to photospheric magnetic field measurements. At chromospheric and photospheric levels the bright features associated to the magnetic network do show intensity variations in many  spectral signatures (\\opencite{lit93}; \\opencite{cau00}). However, the interpretation for these fluctuations is still an open problem: they might be the result of magneto-acoustic waves which propagate upward along the magnetic lines of force (\\opencite{has99}) or they could represent an instability of the chromosphere itself, as a response to micro or nano-flares processes (\\opencite{han97}). These fluctuations are usually studied in a {\\it statistical} sense and have never been put into direct relation with the network flares identified at higher atmospheric layers.  In this paper we will discuss the search performed at chromospheric and photospheric levels for flares occurring on the magnetic network. Beside the customary analysis of the intensity and velocity variations, we also investigate the changes of the magnetic flux density, and their relation with the fluctuations of the velocity field. We use a subset of the high spatial resolution observations obtained in August 1996 during a coordinated observing program between ground based observatory (``R.B. Dunn'' Solar Telescope at NSO/Sacramento Peak) and the Solar and Heliospheric Observatory (SOHO). ",
        "conclusions": "We searched for the presence of chromospheric and photospheric signatures of low-energy flares occurring on the magnetic network. Eleven NBPs, each one corresponding to a patch of definite polarity, were selected in a 2$^\\prime\\times 2^\\prime$ field of view, where also a small and low-activity region was present. The intensity fluctuations of several spectral signatures observed for the NBPs do not show any impulsive episodes typical of the flare occurrence. The general characteristics found for these NBPs do not differ from the ones derived in quiet regions (\\opencite{cau00}). Only for one NBP we found two peculiar characteristics: \\begin{itemize} \\item[--] The intensity fluctuations in the \\Ha~wings ($\\pm$ 1.5 \\AA~) are not in phase for about 20 minutes. From the difference of the intensity of the \\Ha~wings we derived the corresponding velocity and found that three spikes of downward motions (respectively of 2.0, 2.7 and 3.7 $\\pm$ 1.0 Km/s) occurred within this interval. Plasma motions of this kind are generally  considered as a signature of sudden compression of the lower atmosphere, typical of flares.  \\item[--] In the Fourier domain, at the frequency corresponding to 5 minutes oscillations, the MDI magnetic and velocity fluctuations in this NBP show a coherence spectrum with a high value ($\\approx 0.8$, see Figure \\ref{phas_cohe}) when compared to the average for all the other NBPs ($\\approx 0.4$). Analyzing the coherence spectrum for 3 separate periods, we found that before and after the downward velocity episodes, the coherence has an even higher value of about 0.95, while just during the sudden compression episodes, the coherence is lowered to 0.7. We suggest that a reconnection process, that might be responsible for the sudden compression indicated by the three spikes of downward velocity, might also destroy the coherence of magneto-acoustic waves propagating along the magnetic field lines of a long loop connecting this NBP with a far opposite polarity. \\end{itemize}  GOES recorded an A6 burst in both channels, that temporally coincides with the stronger downward velocity spike. We derived  $T_e = 4.5$~ MK and EM $ \\approx 10^{47}$ cm$^{-3}$ for a coronal loop with a constant diameter of 5$^{\\prime\\prime}$ and  a total length of about $8\\times10^4$ Km.  The general picture emerging from the data is hence the following: the particular network bright point in exam is connected via a long magnetic loop to a far opposite polarity structure. In ``quiet'' conditions, magneto-acoustic waves can propagate along this loop as derived from the $-90^{\\circ}$ phase difference measured between B and V fluctuations. Magnetic reconnection involving this loop disrupts the coherence between these fluctuations, and produces the compression episodes observed in the lower atmosphere. At coronal levels the simultaneous increase in SXR emission might identify the corresponding chromospheric evaporation. All this evidence points toward the occurrence of a real ``network flare''. Contrary to more energetic flares, for such an event a sudden intensity increase does not represent a characteristic signature in the low atmosphere. Instead, downward motions seem to be a common property for both larger and smaller flares. An intriguing characteristic of this particular flare is given by the disruption of the coherence between the B and V fluctuations. This is the first time that such a property is observed, so it would be premature to comment on its relevance as a flare signature. Further investigation on the subject is desirable."
    },
    "0011/astro-ph0011119_arXiv.txt": {
        "abstract": "{ We present the results of an analysis of broad band UV observations of the central regions of Abell 1795 observed with the optical monitor on XMM-Newton. As have been found with other UV observations of the central regions of clusters of galaxies, we find evidence for star formation.  However, we also find evidence for absorption in the cD galaxy on a more extended scale than has been seen with optical imaging.  We also report the first UV observation of part of the filamentary structure seen in H$\\alpha$, X-rays and very deep U band imaging.  The part of the filament we see is very blue with UV colours consistent with a very early (O/B) stellar population.  This is the first direct evidence of a dominant population of early type stars at the centre of Abell 1795 and implies very recent star formation at the centre of this cluster. ",
        "introduction": "The final fate of material in the centre of cooling flow clusters remains a mystery.  Perhaps the most obvious end point for matter cooling out of the cooling flow is in stars, and it is indeed true that the cD galaxies in cooling flow clusters often have anomalously blue optical colours (e.g. Allen 1995), implying higher rates of star formation than normal.  There is also a reported correlation between the amplitude and radial extent of the colour anomalies and the inferred mass accretion rate from X-ray measurements (McNamara \\& O'Connell 1992).  However, much of the observations of star formation in clusters of galaxies is based on optical measurements, and so crucial information about the emission from early type, and therefore young, stars is missing.  Unfortunately, UV observations of the cores of cooling flow clusters, where the emission from early type stars would be strongest, are relatively sparse (e.g. A1795 Smith et al. 1997, CL 0939+4713 Buson et al. 1998).  The cD galaxy in Abell 1795 has been extensively studied and shows a number of remarkable features, including extended nuclear emission line gas (Cowie et al 1983; Hu, Cowie \\& Wang 1985; Heckman et al. 1989), an unusual blue continuum (Hu 1992; Allen 1995) and filamentary structure in X-rays (Fabian et al. 2000). These properties have made it one of the best candidates for star formation arising from a cooling flow.  Previous UV studies with UIT (Smith et al. 1997) have shown the presence of emission from early type stars in the cD galaxy with data that were consistent with either continuous star formation or a recent (4 Myr) burst of star formation.  Here we report on new multi-filter UV observations of the central 8 arcminutes of the cooling flow cluster Abell 1795 taken as part of an XMM-Newton PV observation.  Other aspects of the XMM-Newton observations will be discussed elsewhere (Tamura et al. 2001, Arnaud et al. 2001) ",
        "conclusions": "The XMM-Newton OM data have illustrated the usefulness of UV observations in studying populations of early type stars in clusters of galaxies.  We have set constraints on models of star formation at the centre of Abell 1795 and have also detected the presence of extra extinction in the cD galaxy.  We have also presented the first UV observations of emission from the filament seen in deep U band images and in H$\\alpha$ to the south of the cD galaxy.  In particular the observations of the central regions of Abell 1795 have highlighted a population of early type stars located at or near the centre of the gravitational potential that is most likely the consequence of recent star formation.  While the exact significance of this in terms of detailed models of material cooling out of the intercluster medium is not yet clear, these data should, in combination with optical/IR data, provide constraints on the timescale and rate of star formation from the intercluster medium itself."
    },
    "0011/astro-ph0011575_arXiv.txt": {
        "abstract": "We present the DIRTY radiative transfer model in this paper and a companion paper.  This model computes the polarized radiative transfer of photons from arbitrary distributions of stars through arbitrary distributions of dust using Monte Carlo techniques.  The dust re-emission is done self-consistently with the dust absorption and scattering and includes all three important emission paths: equilibrium thermal emission, non-equilibrium thermal emission, and the aromatic features emission.  The algorithm used for the radiative transfer allows for the efficient computation of the appearance of a model system as seen from any viewing direction.  We present a simple method for computing an upper limit on the output quantity uncertainties for Monte Carlo radiative transfer models which use the weighted photon approach. ",
        "introduction": "Dust radiative transfer models are crucial to interpreting many astronomical observations ranging from reflection nebulae with a single star to galaxies comprised of stars, gas, and dust arranged in complex geometries.  Almost all observations are affected by the presence of dust intrinsic to the system under study or present in the foreground.  For dust present in the foreground as is the case for observations of single stars, the effects of dust can be accurately taken into account using a screen geometry.  In this geometry, the effects are well characterized by an extinction curve which is determined by the amount and intrinsic properties (e.g., sizes and materials) of the dust \\citep{cla00}.  However, when the dust is near or mixed with the system being studied, its effects become dependent not only on the intrinsic properties of the dust grains, but on the geometry of the stars and dust in the system as well \\citep{wit92, wit96, wit00}.  Since the effects of mixing stars and dust are non-trivial, a radiative transfer model is required.  Using such a model produces a deeper understanding of the physical properties of the system being studied.  This applies to a wide variety of astrophysical systems, from reflection nebulae to galaxies.  There are a variety of numerical methods other than Monte Carlo techniques which can be used to solve the radiative transfer through dust.  One example is the doubling method \\citep{zub00}.  While the other methods can achieve faster computations for specific cases, Monte Carlo techniques are the most flexible and allow the study of a wider range of astrophysical systems with arbitrary geometries.  There exist many Monte Carlo dust radiative transfer models which have been constructed to study a variety of problems.  In most models, the computational needs of the Monte Carlo technique have been reduced by exploiting symmetries and/or making approximations.  This necessarily restricts the variety of systems to which a particular model can be successfully applied.  Reflection nebulae and circumstellar shells have been successfully approximated as spherically and azimuthally symmetric systems \\citep{wit77, wit82, vos94, cal95, wit96}.  Models for bipolar nebulae, associated with young stellar objects and active galactic nuclei, usually exploit the azimuthal symmetry of such systems \\citep{yus84, whi92, fis94, wod98, you00}.  The general behavior of dust in galaxies has been studied using the concept of spherical galactic environments \\citep{wit92, wit00}.  Spiral galaxies have been studied with models using the approximation of disk galaxies \\citep{bia96, wod97, kuc98}.  The importance of the multi-phase nature of dust distributions has motivated the construction of models which can handle arbitrary distributions of dust.  Models which can handle such dust distributions include those of \\citet{wol98}, \\citet{var99}, \\citet{wod99}, \\citet{bia00}, and our model.  More recently, the importance of the re-emission of the energy absorbed by dust has been acknowledged and increasingly added to some radiative transfer models. The dust re-emission can be broken into three components; thermal equilibrium emission (large particles), non-equilibrium thermal emission (small particles), and the aromatic features (PAH-like particle emission between 3.3 and $\\sim$20~$\\micron$).  The inclusion of the dust re-emission to radiative transfer models has been done with a range of approximations.  The easiest approximation is to include dust in thermal equilibrium \\citep{var99, wol99, bia00b}. Including the non-equilibrium emission and aromatic feature emission is much harder and, as a result, requires a significant amount of CPU time.  Due to the complexity of the non-equilibrium and aromatic feature emission, some models use a simplified \\citep{sil98} and/or empirical approach to including them.  The model described in this paper and a companion paper \\citep{mis00} can handle arbitrary dust distributions and fully includes all three dust emission components self-consistently.  The DIRTY model (DustI Radiative Transfer, Yeah!) grew out of spherically symmetric models which were used to study reflection nebulae \\citep{wit82,gor94,cal95}.  Using Monte Carlo techniques, these models computed the radiative transfer through spherical shells of dust illuminated by point sources.  The model used by \\citet{gor94} allowed for multiple illuminating stars arbitrarily distributed throughout a spherical nebula and the images of such a system could be constructed for arbitrary lines-of-sight.  As such, the \\citet{gor94} model is the direct ancestor of the DIRTY model.  It is important to note that the algorithm used by \\citet{gor94}, while giving the correct total scattered light value, gave too flat a distribution of scattered light across the system.  The correct algorithm for efficient computation of the appearance of a system for a particular line-of-sight is described by \\citet{yus84}.  The motivation for the DIRTY model grew out of the conference ``Dust Survival in Interstellar/Intergalactic Media'' held at the Space Telescope Science Institute in 1994.  One of the major concerns at this conference was that all Monte Carlo radiative transfer models at that time assumed a smooth distribution of dust, yet the interstellar medium is known to be clumpy over a large range of size scales \\citep{col80, sca90, ros95}.  We constructed the DIRTY model to answer this concern.  The main change was to move from spherical shells to rectangular cells as the basic unit of the dust density distribution. Using the simple system of a spherical nebula with a central illuminating star, \\citet{wit96} explored the differences between the radiative transfer in smooth and clumpy 3-dimensional dust distributions.  The implications of clumpy dust for galaxies were discussed by \\citet{wit00} using the concept of spherical galactic environments. The DIRTY model has been applied to starburst galaxies \\citep{gor97, gor00}, spiral galaxies \\citep{kuc98}, the UW Cen reflection nebula \\citep{cla99}, and the nucleus of M33 \\citep{gor99}.  The DIRTY model was constructed to compute the radiative transfer of photons from arbitrary distributions of photon emitters through arbitrary distributions of dust.  This model self-consistently models the scattering and absorption of photons through dust and the re-emission of photons from dust.  The algorithm for the radiative transfer, including polarization, is based on the work of \\citet{wit77}, \\citet{yus84}, and \\citet{cod95}.  The details of this algorithm are the subject of this paper.  The re-emission of energy absorbed by the dust is the subject of a companion paper \\citep{mis00}.  The radiative transfer of the re-emitted photons is handled by the same algorithm as the stellar photons.  This requires an iterative procedure to account for dust absorption and scattering of re-emitted photons and their subsequent re-emission.  The details of the algorithms for the dust emission and iterative procedure are given by \\citet{mis00}. ",
        "conclusions": "We have presented the DIRTY radiative transfer model in this paper and a companion paper \\citep{mis00}.  This model uses Monte Carlo techniques to compute the radiative transfer of photons emitted from arbitrary stellar distributions through arbitrary dust distributions. The weighting scheme used in the DIRTY model results in the efficient computation of the appearance of such systems from any viewing direction.  The dust re-emission includes equilibrium thermal emission, non-equilibrium thermal emission, and the aromatic feature emission.  The dust emission is computed self-consistently with the dust absorption and scattering using an iterative technique.  The details of the dust emission are described by \\citet{mis00}.  The DIRTY model can be used to study a wide range of astrophysical objects due to its flexibility.  The general behavior of multi-phase dust has been studied using the DIRTY model in reflection nebula environments \\citep{wit96} and galactic environments \\citep{wit00}. The very non-symmetric reflection nebula surrounding the R CrB star UW Cen was successfully modeled using the DIRTY model \\citep{cla99}.  The optical depth of spiral galaxies was investigated using an early version of the DIRTY model \\citep{kuc98}.  The finding that the lack of a 2175~\\AA\\ depression in the spectra of UV-selected starburst galaxies was not due to radiative transfer effects but due to the dust grain properties was accomplished with the DIRTY model \\citep{gor97}. The ultraviolet through near-IR spectral energy distribution of the M33 nucleus was modeled using the DIRTY model and discovered to be a post-starburst stellar population surrounded by a shell of Milky Way-like dust \\citep{gor99}.  The general behavior of the infrared emission of starburst galaxies has been studied by \\citet{mis00}.  The full ultraviolet through far-infrared spectral energy distributions of a handful of starburst galaxies was studied by \\citet{mis00b}.  The range of objects which the DIRTY model has already been applied to is a good illustration of the usefulness of such a radiative transfer model.  In addition to continuing to use the DIRTY model to explore the nature of dust and its affects on a variety of astrophysical systems, we plan to continue to improve the model itself.  One such improvement would be to modify the current way the dust distribution is stored to allow for grid points to be subdivided independently allowing for a larger range of size scale to be efficiently modeled \\citep{wol99}.  For example, this would allow for modeling the physically small region near the accretion disk in a QSO while simultaneously modeling the large scattering regions in the jets of the QSO.  Another improvement would be to include electron scattering which would improve the application of the DIRTY model to the study of QSOs.  Finally, the code could be parallelized to allow larger dust distribution grids to be studied in shorter amounts of time.  As there are always more interesting astrophysical systems to be studied than time in the day, anyone interested in using the DIRTY model should contact the authors for possible collaboration."
    },
    "0011/astro-ph0011096_arXiv.txt": {
        "abstract": "The unusual morphology of LRG J0239$-$0134 at $z$ =1.062 has been interpreted as a ring galaxy. Here we propose an alternative idea that the ring-like morphology is attributed to a pair of superbubbles driven by the intense starburst in the central region of this galaxy. Supporting evidence for a superbubble model is; 1) the poststarburst nature in the central body suggests that a burst of supernova explosions could occur at least $\\sim 10^7$ yr ago, 2) the dark lane seen in the central body suggests that we observe this object from a nearly edge-on view, and, 3) the ring-like morphology is not inconsistent with an idea that it is a pair of superbubbles. All these pieces of evidence for the superbubble model seem circumstantial. However, if this is the case, this galaxy provides an important example of the superwind activity at high redshift. ",
        "introduction": "The greatest observational capability of the {\\it Hubble Space Telescope (HST)} has been providing us a lot of invaluable information on the evolution of galaxies (e.g., Williams et al. 1996; Lanzetta, Fern\\'andez-Soto, \\& Yahil 1996; van den Bergh et al. 1996; Cowie, Hu, \\& Songaila 1995; van Dokkum et al. 1999). In addition, massive and clusters of galaxies have been also used as natural gravitational telescopes to investigate properties of high-$z$ galaxies [Smail et al. 1996; Colley, Tyson, \\& Turner 1996, Franx et al. 1997; Soucail et al. 1999 (hereafter S99); P\\'ello et al. 1999; B\\'ezecourt et al. 1999].  Since optical observations of high-$z$ galaxies probe their redshifted ultraviolet radiation, we must be careful in interpreting the data because active star-forming regions may be selectively shown up, making often galaxy morphology knotty or irregular (Bohlin et al. 1990; Giavalisco et al. 1996; Hibbard \\& Vacca 1997). However, on one hand, it seems likely that young galaxies have more chaotic morphological properties because they are dynamically younger than present-day galaxies. For example, galaxies with chain-like morphology have been found in {\\it HST} deep images (Cowie et al. 1995; van den Bergh et al. 1996; Abraham 1997; Bunker et al. 2000; Taniguchi \\& Shioya 2001). It is also mentioned that more frequent collisions between/among galaxies (or subgalactic building blocks) result in peculiar morphology of high-$z$ galaxies (e.g., van den Bergh et al. 1996).  Among galaxies with peculiar morphologies, one interesting object is a ring-like galaxy at $z = 1.062$ (LRG J0239$-$0134) found as an optical counterpart of the 850 $\\mu$m source, SMM J02399$-$0134 (Smail et al. 1998: S99). Such ring galaxies are not so rare in the local universe and their formation mechanism is now well known; the axial penetration of a smaller companion galaxy to a disk galaxy (Theys \\& Spiegel 1976, 1977; Lynds \\& Toomre 1976; Appleton \\& Marston 1997 and references therein). Therefore, if LRG J0239$-$0134 is a ring galaxy, it would have a less-luminous companion galaxy. Although two galaxies are close to LRG J0239$-$0134 (see Figure 1), one of them, \\#32 in Mellier et al. (1998), belongs to the cluster of galaxies Abell 370 ($z = 0.37$) and the other galaxy A in S99, or L4 in Barger et al. (1999), has a redshift of 0.42. Thus, these two galaxies have no relation to  LRG J0239$-$0134. Although another two very faint objects (B and C in Figure 1) are also seen to the south of LRG J0239$-$0134, they appear too faint to be a colliding partner. The most probable candidate for the colliding partner seems an object seen just to the north of the main  body of LRG J0239$-$0134 (D in Figure 1) although we cannot rule out the possibility that this object might not be separate from the main body of LRG J0239-1034. However, there is no spectroscopic confirmation. Hence it seems still uncertain whether or not LRG J0239$-$0134 is a {\\it usual} ring galaxy made by the axial penetration of a companion galaxy. In this Letter, we propose an alternative idea that the ring-like morphology is attributed to a pair of superbubbles driven by the intense starburst in the central region of this galaxy. ",
        "conclusions": "\\subsection{LRG J0239$-$0134 as an Ultraluminous Infrared Galaxy}  Photometric data from rest-frame ultraviolet (UV) to MIR of LRG J0239$-$0134 were obtained by S99. In addition to these data, Barger et al. (1999) reported its 850 $\\mu$m flux, 11 mJy. Using these data, we investigate the overall spectral energy distribution (SED) of LRG J0239$-$0134. The SED is shown in Figure 3 together with those of NGC 6240 (Griffith et al. 1995; Lisenfeld et al. 1996, 2000; Klaas et al. 1997; Spinoglio et al. 1995; White \\& Becker 1992) and  Arp 220 (Rigopoulou, Lawrence, \\& Rowan-Robinson 1996; Klaas et al. 1997). As shown in Figure 3, the SED of LRG J0239$-$0134 is much similar to that of NGC 6240 than to that of Arp 220. This makes sense because NGC 6240 harbors a luminous active galactic nucleus while Arp 220 may not. Fitting the two SEDs of NGC 6240 and Arp 220 to the observed SED of LRG J0239$-$0134, we obtain the total IR luminosity between 8 $\\mu$m and 1000 $\\mu$m, $L$(8-1000$\\mu$m) $\\sim 7.9 \\times 10^{12} h_{50}^{-2} L_\\odot$ and $8.6 \\times 10^{12} h_{50}^{-2} L_\\odot$, respectively; in the case of $ h_{50}=1.5$, $L$(8-1000$\\mu$m) $\\sim 3.5 \\times 10^{12} L_\\odot$ and $\\sim 3.8 \\times 10^{12} L_\\odot$, respectively.  For nearby starburst galaxies including some local ULIGs, Mouri et al. (1990) found a relation of $L_{\\rm FIR} \\sim 10^3 L_{\\rm 3.3\\mu m}$ where $L_{\\rm FIR}$ is the FIR luminosity derived from rest-frame 60 $\\mu$m and 100 $\\mu$m fluxes (e.g., Sanders et al. 1988; Sanders \\& Mirabel 1996). Using the observed $L_{\\rm 3.3\\mu m}$ (S99), we would obtain $L_{\\rm FIR} \\sim 4 \\times 10^{13} L_\\odot$. However, from the SED fitting, we obtain $L_{\\rm FIR} \\sim 8 \\times 10^{12} h_{50}^{-2} L_\\odot$ at most, being lower by a factor of five  than the above expected value. Although the SED concordance between high-$z$ submm sources and local ULIGs has been claimed (Ivison et al. 1998; see also Trentham 2000), it is suggested that LRG J0239$-$0134 is a different population from local {\\it starburst-dominated} ULIGs though its IR luminosity is comparable to those of the ULIGs\\footnote{ Typical ring galaxies in the local universe have FIR luminosities of $L_{\\rm FIR} \\sim 4 \\times 10^{10} h_{50}^{-2} L_\\odot$ (Appleton \\& Struck-Marcell 1987), being much fainter than that of LRG J0239$-$0134. On the other hand, they have ring radii of 20 - 40 kpc (Appleton \\& Struck-Marcell 1987), being significantly larger than that of LRG J0239$-$0134.}.  \\subsection{Concluding Remarks}  We have presented our superbubble model for the unusual ring-like galaxy LRG J0239$-$0134 at $z = 1.062$. All pieces of supporting evidence for the superbubble model seem circumstantial. However, if this is the case, this galaxy provides an important example of the superwind activity at high redshift.  Similar examples may be  the extended Ly$\\alpha$ blobs (LABs) at $z \\approx 3.1$ found by Steidel et al. (2000) because their observational properties are also interpreted in terms of superwind activities (Taniguchi \\& Shioya 2000; cf. Haiman, Spaans, \\& Quataert 2000; Fardal et al. 2000). Taniguchi \\& Shioya (2000) proposed an evolutionary connection from dust-enshrouded (or dusty) submm sources (hereafter DSSs; e.g., Hughes et al. 1998; Barger et al. 1998; Barger et al. 1999; Ivison et al. 2000 and references therein) to the LABs. It is interesting to note that LRG J0239$-$0134 can be considered as a link between the DSSs and the LABs. UV imaging of LRG J0239$-$0134 will be important to probe its extended Ly$\\alpha$ emission.  Our model predicts that the central body of LRG J0239$-$0134 is rotating along the minor axis. Further, the ring-like regions are affected by shock heating to some extent. High-resolution optical spectroscopy will be useful in examining these predictions.  \\vspace {0.5cm}  We would like to thank Kazushi Iwasawa and Yasuhiro Shioya for useful discussion and comments. We would also like to thank an anonymous referee for useful comments and suggestions. YT thanks Rolf-Peter Kudritzki, Bob McLaren, and Dave Sanders at Institute for Astronomy, University of Hawaii for their warm hospitality. This work was financially supported in part by the Ministry of Education, Science, and Culture, Japan (grants 10044052, and 10304013)."
    },
    "0011/astro-ph0011269_arXiv.txt": {
        "abstract": "We have observed the B\\,{\\sc i} line at 2089.6\\,{\\AA} in the metal poor star HD\\,140283 with HST/STIS. The observation was an attempt to confirm the B abundances derived from the 2497\\,\\AA\\ B\\,{\\sc i} line in earlier works in general, and for for this star in particular (Edvardsson et al. 1994;  Kiselman \\& Carlsson 1996). The resulting spectrum gained from 8 orbits of observations is hardly consistent with the boron abundance derived in the earlier works. A pure $^{10}$B line could produce the feature observed but no reasonable process can produce such a ratio between $^{10}$B and $^{11}$B. More likely the analysed feature is affected by a statistical fluctuation. A conservative upper limit estimate of the NLTE corrected boron abundance gives a value marginally consistent with that derived from the 2497\\,\\AA\\ line. We are proposing further observations to derive a more definite abundance from the line. ",
        "introduction": "The abundances of the lightest elements and their isotopes play a well-known and important r\\^ole in cosmology, for the early evolution of the Galaxy, and for the understanding of stellar evolution. Be and B abundances in the oldest stars have important implications for the spallative interactions between cosmic rays and the ISM which are thought to be mainly responsible for their production. There is, however, still no generally accepted model for the details of these mechanisms, and improved abundance determinations -- especially for the extreme metal poor stars -- are very important for the testing of different scenarios. In conjunction with measurements of Li abundances, Be and B also probe the convection zones and mixing processes in late-type stars.  In halo stars, all boron abundance determinations have so far been derived from the 2496.7\\,\\AA\\ B\\,{\\sc i} line. This is an attempt to improve this state of affairs by observations of the 2089.6\\,\\AA\\ B\\,{\\sc i} line in the well-known, bright halo dwarf star HD\\,140283 (BD\\,-10 4149, HIP\\,76976, SAO\\,159459). This line has the extra virtue of showing an unusually large isotope shift between the two stable isotopes $^{11}$B and $^{10}$B, (25\\,m\\AA, Johansson et al. 1993) which has already been used or tried for isotope-ratio determinations in the ISM and in galactic disk- or thick-disk stars. This will eventually also be possible for halo dwarfs. ",
        "conclusions": "We have used eight HST/STIS orbits to observe the 2089.6\\,\\AA\\ B\\,{\\sc i} line in HD\\,140283. The observed spectrum is only marginally consistent with the boron abundance found in previous determinations based on the 2496.7\\,\\AA\\ B\\,{\\sc i} line. We are proposing further HST/STIS observations to secure the detection of the 2089.6\\,\\AA\\ B\\,{\\sc i} line, as a very important check on all previous boron abundance determinations for halo stars."
    },
    "0011/astro-ph0011025_arXiv.txt": {
        "abstract": "Prospects for testing the adiabaticity of the primordial cosmological perturbations using MAP and PLANCK are evaluated. The most general cosmological perturbation in a universe with just baryons, photons, neutrinos, and a cold dark matter (CDM) component is described. In addition to the familiar adiabatic mode, there are four nonsingular isocurvature modes: a baryon isocurvature mode, a CDM isocurvature mode, a neutrino density isocurvature mode, and a neutrino velocity isocurvature mode. The most general perturbation is described by a $5\\times 5$, positive-definite, symmetric, matrix-valued function of wave number whose off-diagonal elements represent the correlations between the above mentioned modes. We found that when three modes and their correlations are admitted, the fractional uncertainties in the cosmological parameters and amplitudes of the isocurvature modes and their correlations become of order unity. These degeneracies, however, can be broken with the polarization information provided by PLANCK, reducing the uncertainties to below the ten percent level. Polarization is thus crucial to testing the adiabaticity of the primordial fluctuations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011213_arXiv.txt": {
        "abstract": "We have observed in the mid-infrared (4-18\\,$\\mu$m) the counterpart of the compact object GRS\\,1915+105, and the western jet of SS433. The images were carried out with the ISOCAM infrared camera on board of the Infrared Space Observatory (ISO). The mid-infrared photometry of GRS\\,1915+105 shows the presence of an additional contribution besides the synchrotron emission. The 15\\,$\\mu$m images of the large-scale western lobe of the SS433/W50 nebula are compared to the radio and X-ray ones. They show infrared synchrotron emission on the western edge of SS433/W50 lobe. ",
        "introduction": "We have imaged the close environment of GRS\\,1915+105 with ISOCAM at two epochs, on 1996 April 28 and on 1997 October 24. The binary system appears in mid-infrared wavelengths as a faint point source, globally brighter in our 1997 observations than in the 1996 ones. No elongated structure was observed with 1.5$''$ pixel field of view, the best ISOCAM resolution. We obtained the photometry with several large-band filters between 4 and 18\\,$\\mu$m, and have compared it to simultaneous observations at near-infrared, radio and X-ray wavelengths.  On 1996 April 28 GRS\\,1915+105 was in a High/Soft state, with nearly constant RXTE/ASM flux ($\\sim$ 90\\,counts/s), no significant BATSE flux, and $\\sim$\\,3\\,mJy at 15\\,GHz (Ryle Telescope) few days before and after. On 1997 October 24, 5 days before a major radio and X-ray flare, the compact source was in a plateau state with $\\sim$ 50\\,mJy at 2.25 \\& 8\\,GHz (GBI) giving a radio spectral index $\\sim -0.08$, with low constant RXTE/ASM flux ($\\sim$ 20\\,counts/s) and BASTE photon flux $\\sim$ 0.06.  The IR data have been de-reddened according to the law from Lutz et al.\\,(1996). We used very simple models combining a hot black body emission or a hot synchrotron one to fit the near infrared data, with a cold black body emission (for a dust envelope heated by the binary) and a synchrotron emission extrapolated from radio wavelengths. We conclude that synchrotron emission alone is not sufficient to explain the mid-infrared flux of GRS\\,1915+105, and that an additional contribution, which could be dust emission around the binary, is needed. ",
        "conclusions": ""
    },
    "0011/astro-ph0011449_arXiv.txt": {
        "abstract": "Oxygen is a much better evolutionary index than iron to follow the history of Lithium-Beryllium-Boron (LiBeB) since it is the main producer of these light elements at least in the early Galaxy. The O-Fe relation is crucial to the determination of the exact physical process responsible for the LiBeB production. Calculated nucleosynthetic yields of massive stars, estimates of the energy cost of Be production, and above all recent observations reported in this meeting seem to favor a mechanism in which fast nuclei enriched into He, C and O arising from supernovae are accelerated in superbubbles and fragment on H and He in the interstellar medium. ",
        "introduction": "Lithium-Beryllium-Boron take a special status in the general framework of nucleosynthesis. These nuclei, are indeed of exceptional fragility and they are destroyed in stars above about 1 million degrees. The main formation process available is spallation, i.e. fragmentation of medium light isotopes (CNO) leading to lighter species as ${^6}Li$, ${^7}Li$, ${^9}Be$, $^{10}B$ and $^{11}B$. The physical parameters  of the spallation mechanism are fourfold: i) the production cross sections as a function of energy  are fairly well measured in the laboratory ii) the source composition of the energetic component ii) the associated (injection) energy spectrum iv) the target composition. In the following, we describe the two reverse spallative processes able to produce LiBeB and we confront them to observational constraints. We show how the relation between oxygen and iron, in the halo phase is determining to discriminate between the two processes. ",
        "conclusions": "The synthesis of LiBeB in the halo proceeds through nuclear spallation, essentially by the break up of oxygen and the $\\alpha + \\alpha$ reaction. The observed relation between BeB and Fe has to be translated through the O - Fe relation into a BeB - O which is representative of the relevant physical production process. And then the observational O - Fe relation is determining. (except if Be, B and O are measured simultaneoulsy in stars  (A.M. Boesgaard and K. Cunha, this meeting). Anyway the primary component is required in the halo phase, then afterwards the secondary process takes over; the question is therefore at which metallicity the transition occurs, the answer depends of the behaviour of [O/Fe]. The ideal test of the scenarios would be to get the evaluation of O, Fe, Mg Be, B, Li abundances in the same stars. Impressive progress are waited from the VLT. Gamma ray line astronomy, through the european INTEGRAL satellite, to be lauched in 2002, will help to constrain the energy spectrum and intensity of LEC."
    },
    "0011/astro-ph0011163_arXiv.txt": {
        "abstract": "\\ngc\\ is an optical nebula excited by the powerful wind and the radiation of the Wolf--Rayet star \\hd. We have investigated the interaction between this massive star and the surrounding neutral gas by analyzing the large-scale 21cm-\\hi\\ emission and by mapping the nebula in the J = \\low\\ and the J = \\hig\\ lines of CO. We found a conspicuous ($70 \\times 37$ pc) \\hi\\ shell, expanding at 12 \\kms, likely produced during the main-sequence phase of the star. The molecular gas towards \\ngc\\ shows three velocity components. Two of these components, A1 and A2, have narrow linewidths (1--2 \\kms) and radial velocities of 35--38 and 64--68 \\kms, respectively. The third component is detected at radial velocities between 50 and 58 \\kms\\ and has a broader profile (up to 5.5 \\kms). Furthermore, this component is morphologicaly related with the nebula and has a velocity gradient of a few \\kms. We have also estimated the physical parameters of the molecular gas by means of a LVG modelling of the CO emission. The gas projected onto the southern \\hii\\ region of the nebula has low CO column density and is rather hot, probably up to 80\\,K. Several profiles of the $^{13}$CO J = \\low\\ line near the peak of the emission, together with a weak emission bridge between the broad and one of the narrow components (component A2), suggest the presence of a shock front acting in the southern part of the nebula. This shock was likely produced in a previous RSG/LBV phase of \\hd. ",
        "introduction": "The different evolutive stages of massive stars have strong influence on their surroundings and, consequently, on the galactic evolution. During their lives, the massive stars inject large amounts of matter, energy and momentum into the interstellar medium (ISM). In the pre-main sequence, massive stars undergo a phase of very energetic mass loss giving rise to molecular outflows. UV radiation, expansion of \\hii\\ regions and fast stellar winds are present in the main sequence stage. In the short--lived red supergiant (RSG) and luminous blue variable (LBV) stages, the winds become slower and denser and ejections of up to some solar masses may be produced. The Wolf--Rayet (WR) stage, prior to SN explosions, is characterized by a copious mass loss (typically $10^{-5}$ \\mlr) driven by a fast (about 2000 \\kms) and chemically enriched stellar wind. Massive stars are therefore continuously heating, ionizing, shocking and blowing-up the ISM surrounding them. Garc\\'{\\i}a-Segura \\& Mac Low (1995) have modeled the evolution of the gas which surrounds massive stars, from the main sequence to the WR stage. They only considered simple models for the mass ejection in the form of stellar winds in the different stages of evolution and predict the presence of large scale structures around the massive stars. Both the blown-up gas and the ejecta form multiple shells as a consequence of the different shockfronts produced inside the gas.  During the main sequence phase, the gas around massive stars is ionized, heated and highly evacuated by the fast, long-lived stellar wind and the Lyman continuum radiation. A hot cavity is surrounded by an expanding shell. The ionization front may get trapped within it, the shell may start to recombine and become observable in the line radiation emitted by neutral atoms or molecules. These large bubbles have been successfully detected in the 21-cm line of atomic hydrogen (\\cite{arn92}; \\cite{riz94}; \\cite{cap98}; \\cite{riz98}; \\cite{arn99}), in the form of huge shells of 30--70 pc, expanding at 6--20 \\kms, with masses of hundreds (or even thousands) of solar masses. Some of the large IRAS shells found by Marston (1996) might also be related to this type of structures.  Once the massive star finishes the hydrogen burning in its nucleus and becomes a RSG or LBV, its stellar wind suddenly becomes more dense, and a significant amount of the stellar mass is deposited into the ISM at velocities around 20 \\kms. Very probably, the formation of the optical ring nebulae starts in this brief stage. When the star comes to a WR stage, the wind is accelerated to hypersonic velocities and rapidly reaches the RSG/LBV wind. Many of the optical nebulae possess arcs and have high inhomogeneities in the abundances of heavy elements (\\cite{chu83}, 1999).  However, the effects of the WR phase in the neutral (neither atomic nor molecular) gas have not been easily detected. The structures formed during this phase would not be as large as those developed in the O-phase. Just in few cases, the neutral atomic hydrogen has been detected in interaction with ring nebulae (\\cite{cap96}; \\cite{arn96}). Vibrationally excited molecular hydrogen has only been detected towards \\ngc. The CO J = \\low\\ emission has been observed around WR\\,16 (\\cite{mar99}) and in \\ngc\\ (\\cite{sch81}).  \\ngc\\ is a very interesting object, and many aspects of which has been thoroughly studied in the recent years. This optical nebula is excited by \\hd\\ (WR\\,7 in the catalogue of van der Hutch et al.\\ (1988), a WN4 star located in the outer Milky Way. Its distance varies from 3.5 to 6.9 kpc depending on the authors (see \\cite{gou94} for a discussion). The mass--loss rate ($<7\\times 10^{-5}$ \\mlr) and the terminal wind velocity (1545 \\kms) have been measured by Schneps et al.\\ (1981) and Rochowicz \\& Niedzielski (1995), respectively. \\ngc, a prototypical wind-blown bubble (\\cite{chu83}), is nearly spherical with several small filaments inside. The nebula is not highly enriched of heavy elements as a whole, with the exception of the filaments (\\cite{est90}). This may indicate that the nebula was largely produced in the RSG/LBV and the WR stages of the exciting star. From the optical emission lines (\\cite{gou83}) and radio recombination lines (\\cite{loc89}; \\cite{fic91}), we know that the ionized gas has several components at velocities between 30 and 70 \\kms\\ (all velocities in this paper are referred to LSR).  An \\hii\\ region partially surrounds the wind-blown bubble and is  apparently limited by obscuring material. Schneps et al.\\ (1981) have taken several profiles in the J = \\low\\ line of CO and detected the presence of three  clouds at 37, 54 and 67 \\kms. The presence of broad CO profiles at 54 \\kms\\ and  the morphology of the emission, just following the south-eastern edge of the nebula, suggests that this component has been interacting with the nebula. St-Louis et al.\\ (1998) have analyzed the \\htwo\\ emission in this region. They have found the 1--0 S(1) line towards the southern border of the \\hii\\ region, but they could not establish the nature of the excitation (fluorescence or  shocks) of the \\htwo. Cappa et al.\\ (1999) made a complete map of the ionized component at the 1465 MHz continuum, and traced the \\hi\\ emission presumably connected with the nebula. These \\hi\\ features are probably associated with the WR stage of \\hd; so far, the large atomic main-sequence bubble predicted by the theory has not been yet reported.  The goal of this paper is to look into the history of the interaction of the wind and the radiation of \\hd\\ and its surroundings. We have analyzed a public \\hi\\ survey and studied for the first time the large scale interaction of \\hd\\ in the O-phase with the ambient ISM. Based on the mapping of the J = \\low\\ and J = \\hig\\ lines of CO over the whole optical nebula, we study in detail the morphology and kinematics of the CO around \\ngc. We also present several $^{13}$CO profiles towards selected positions in order to determine the global physical parameters of the molecular gas, such as kinetic temperature, density and  opacity of the lines. This study has allowed us to shed some light on the dominant excitation mechanism of the neutral gas in the WR stage. ",
        "conclusions": "We have analyzed the large-scale 21 cm-\\hi\\ emission in a $6\\degr \\times 6\\degr$ field around the nebula \\ngc, and found a large expanding bubble roughly centered at the nebula and its exciting star, the Wolf-Rayet \\hd. The systemic velocity of the bubble is $\\sim$ 64 \\kms\\ and it is expanding at 12 \\kms. At an assumed distance of 5 kpc, the bubble has a size of $70 \\times 37$ pc and a \\hi\\ mass between 700 and 2400 \\msun. By simple estimates of energetics, ages and stellar wind parameters, we think that this feature is a wind-blown bubble produced by the O-progenitor of the WR star. The presence of a nearly spatially-coincident IRAS shell reinforces our hypothesis.  We have also mapped the CO J = \\low\\ and \\hig\\ emission adjacent to \\ngc\\ and found three CO components in this region. Two of these components, with radial velocities of 34 (component A1) and 67 (component A2) \\kms, have narrow profiles (up to 2 \\kms) and do not show any morphological or kinematical effects which indicate the disturbance by the nebula or the WR star. However, the third component (the component B, emitting at 54 \\kms), clearly bounds \\ngc\\ at its southern and eastern border. The profiles are broad, with linewidths of 4--5.5 \\kms. A velocity gradient of a few \\kms\\ is noted towards the south and south-eastern interface with the nebula. A weak ``bridge'' in velocity of small angular extension seems to connect this component with the component B2 at the southern part.  We have observed the $^{13}$CO J = \\low\\ line at a few points in the region where the component B is detected. We have found a striking variation of the $^{12}$CO/$^{13}$CO ratio both in position and in velocity. This ratio seems to decrease near the peak of the CO emission, for velocities higher than 53 \\kms. This particular emiting region is not resolved by the beam, and is probably 0.2 pc wide.  From the observational data, we have estimated the physical conditions for each CO component using the LVG approximation. The highest excitation temperature and density are found in the component B. It is remarkable that high temperatures (up to 80\\,K) are derived for the CO emission projected onto the \\hii\\ region .  In view of the kinematics, the morphology and the physical properties of the CO, we think that the component A2 was shocked and accelerated by the expanding bubble to reach the radial velocities of the component B. The shock front is still acting at the southern part of the region. The origin of such shock might be related to the WR phase of \\hd\\ or, more probably, to previous episodes of RSG of LBV. From the $^{12}$CO/$^{13}$CO ratio we infer the presence of a thin layer of dense gas which represent the more recently shocked material. Our data, however, cannot resolve spatially the region where a thin layer of shocked material is expected. CO observations with better resolution and sensitivity is needed to disclose the main features of the region."
    },
    "0011/astro-ph0011355_arXiv.txt": {
        "abstract": "The CANGAROO-III project, a system of four 10~m telecopes dedicated for gamma-ray astrophysics, has started in 1999 and is expected to complete in 2004 in Woomera, South Australia. We report the construction work during the first year, which includes the completion of the first 10~m telescope built as a 7~m telescope in 1999, and the work in progess to increase the performance by constructing three more similar telescopes based on the experience gained during the construction of the first telescope. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011388_arXiv.txt": {
        "abstract": "We review current and new estimates of the eclipsing binary YZ Cassiopeiae metallicity (Z).  Since the individual components cover a quite large range of mass (1.35-2.31M$_{\\odot}$), YZ Cas is potentially one of the best stellar laboratories to understand the structure and evolution of 1 to 2 M$_{\\odot}$ stars. The derivation of Z from IUE spectra, as well as from photometric indices, provides the chemical composition of the atmosphere (Z$_{\\rm atmospheric}$), while the fit of evolutionary tracks provides the initial chemical composition (Z$_{\\rm initial}$). While a disagreement is expected between Z$_{\\rm atmospheric}$ and Z$_{\\rm initial}$ because the primary component is an Am star (one expects Z$_{\\rm atmospheric}$ to be larger), we find some unexpected discrepancy between atmospheric determinations of Z for this star. ",
        "introduction": "$ from IUE spectra}  From IUE spectra of the YZ Cas primary, [LM83] found that the metals add up to a Z value of over 0.03 (Z$=$0.036$\\pm$0.005). We try to assess to what extent this high Z-value has to be revised. It is well known that the abundance of iron is enhanced in Am stars with respect to normal  A-F stars. [LM83] found for YZ Cas A that Fe is $\\sim$6  times overabundant, which implies that Fe contributes to $\\sim$43\\% of Z. In order to update the [LM83] study, we applied two corrections to this value: 1) They assumed an iron solar abundance A$_{Fe}$$=$7.60. More detailed studies give A$_{Fe}$ $=$ 7.50$\\pm$0.05 ([GS99]) implying that Fe contributes to $\\sim$57\\% of Z. 2) Their analysis assumed T$_{\\rm eff}$$=$10,300K. Recent determinations give 9100$\\pm$300K ([R00]). This would reduce the [LM83]'s iron abundance by about 0.15-0.20 dex, and then imply that Fe contributes to $\\sim$33\\% of Z. Both corrections are opposite and imply altogether that Fe contributes to $\\sim$40\\% of Z. Hence, our attempt to update some of the [LM83] assumptions does not change their high-Z conclusion. ",
        "conclusions": ""
    },
    "0011/astro-ph0011341_arXiv.txt": {
        "abstract": "We present synthesis models of the X--ray background where the available X--ray observational constraints are used to derive information on the AGN population properties. We show the need for luminous X--ray absorbed AGNs, the QSO2s, in reproducing the 2--10 keV source counts at relatively bright fluxes. We compare a model where the evolution of absorbed AGNs is faster than that of unabsorbed ones, with a standard model where absorbed and unabsorbed AGNs evolve at the same rate. It is found that an increase by a factor of $\\sim 2$ from $z=0$ to $z\\sim1.3$ in the ratio between absorbed and unabsorbed AGNs would provide a significant improvement in the data description. Finally, we make predictions on the AGNs to be observed in deep X--ray surveys which contain information on the AGN space density at high redshift. ",
        "introduction": "Recent surveys have provided strong evidence that the extragalactic X--ray background (XRB) is mainly due to the integrated emission of single sources. In the soft 0.5--2 keV energy range about 70--80\\% of the XRB intensity has been resolved by ROSAT at a limiting flux of $10^{-15}$ erg cm$^{-2}$ s$^{-1}$ (Hasinger et al. \\cite{hasinger}). The resolved fraction increases up to $\\sim 90\\%$ when adding the Chandra results down to a flux of $2.3\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$ (Mushotzky et al. \\cite{musho} ; Giacconi et al. \\cite{giacconi}). In the hard 2--10 keV energy range ASCA and BeppoSAX resolved about 25--30\\% of the XRB (Ueda et al. \\cite{ueda}; Della Ceca et al. \\cite{rdc2}; Giommi, Perri \\& Fiore \\cite{giommi}). When adding Chandra data (Mushotzky et al. \\cite{musho}; Giacconi et al. \\cite{giacconi}) more than 60\\% of the hard XRB is resolved. Most of the ROSAT and ASCA sources have been optically identified as Active Galactic Nuclei (AGNs; Schmidt et al. \\cite{schmidt2}; Akiyama et al. \\cite{ak00}, hereafter Ak00). The optical identifications for the Chandra sources are not complete, however AGNs seem to constitute the main population (Barger et al. \\cite{barger}). These results have confirmed the predictions of AGN synthesis models of the XRB, originally proposed by Setti \\& Woltjer (\\cite{setti}) and subsequently refined into more detailed variants (e.g. Madau, Ghisellini \\& Fabian \\cite{madau}; Comastri et al. \\cite{comastri}); the XRB spectrum is reproduced by summing the contribution of unobscured and obscured AGNs. In these models, the hard spectrum of obscured AGNs is crucial in reproducing the XRB spectral shape above 2 keV.  The existence of a population of luminous absorbed AGNs, the QSO2s, is still a debated issue. Some QSO2s could be hidden in the Ultraluminous IRAS Galaxies (ULIRGs; Kim \\& Sanders \\cite{kim}). Indeed, many ULIRGs do show AGN signatures in their spectra (Veilleux, Kim \\& Sanders \\cite{veilleux}). Also, examples of luminous absorbed AGNs could be represented by ROSAT sources with very red colours ($R-K\\ga5$)  detected at faint fluxes in the Lockman Hole (Hasinger et al. \\cite{hasinger2}). Finally, another interesting possibility, which however is still to be verified, is that QSO2s appear as normal blue, broad lined QSOs in the optical. While luminous AGNs with only narrow optical lines are not commonly observed (see e.g. Halpern, Turner \\& George \\cite{halpern}), ASCA and BeppoSAX have found luminous AGNs with a hard X--ray spectrum showing broad lines in the optical (Ak00; Fiore et al. \\cite{fiore}). If the X--ray hardness is due to cold gas absorption, one could therefore suspect that these objects have different dust--to--gas properties with respect to low luminosity AGNs, where X--ray absorption is generally coupled with an optical type 2 (only narrow lines) classification (Maiolino et al. \\cite{maio3}; Risaliti et al. \\cite{risaliti2}).  Recent works (Gilli et al. \\cite{gilli}, hereafter Paper I; Pompilio et al. \\cite{pompilio}) suggested that the evolution of obscured AGNs might be different from that of unobscured ones, with the number ratio R between obscured and unobscured AGNs increasing with redshift. Observational data seem to confirm this suggestion. Indeed, Reeves \\& Turner (\\cite{reeves}) found that in a sample of radio loud and radio quiet QSOs observed by ASCA the fraction of X--ray absorbed AGNs increases with redshift.  The density of high redshift AGNs has been determined in the optical (Schmidt, Schneider \\& Gunn \\cite{schmidt}) and in the radio (Shaver et al. \\cite{shaver}) bands. It is found that the AGN density declines above $z\\sim3$, as recently confirmed by results from the Sloan Digital Sky Survey (Fan et al. \\cite{fan}). In the X--rays the situation is uncertain. The behaviour of the soft X--ray Luminosity Function (XLF) derived by Miyaji, Hasinger \\& Schmidt (\\cite{miyaji}, hereafter Mi00a) from ROSAT data is indeed consistent with a constant AGN density above $z\\sim2$.  In this work we examine XRB synthesis models under different assumptions in the input parameters. We check the differences in reproducing all the available X--ray constraints between a standard model where the ratio R does not evolve with the redshift and a model where absorbed AGNs evolve more rapidly than unabsorbed ones. We analyze the effects of neglecting QSO2s when the XRB intensity is reproduced only with unobscured AGNs and low luminosity absorbed AGNs. Finally, we compare two alternative scenarios where the density of AGNs at high redshift is constant or declines above $z\\sim3$.  Throughout this paper the deceleration parameter and the Hubble constant are given the values $q_{0}=0.5$ and $H_{0}=50$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have presented different synthesis models for the XRB, checking the model predictions with all the available X--ray constraints posed by the source counts, XLF, redshift distributions and absorption distributions. Models assuming a population of high luminosity obscured AGNs, the QSO2s, are found to be consistent with all the available data. Furthermore, the existence of QSO2s is found to be necessary in reproducing the 2--10 keV source counts at relatively bright fluxes ($\\sim 10^{-13}$ erg cm$^{-2}$ s$^{-1}$). We found that a model (model B) where the evolution of absorbed AGNs is faster than that of unabsorbed ones provides a better description of the data with respect to a standard model (model A) where absorbed and unabsorbed AGNs evolve with the same rate. Our models are also in agreement with the first results from Chandra deep surveys. Chandra sources, together with those detected in XMM deep surveys, will provide additional informations on the QSO2 population and will verify if the faster evolution of absorbed AGNs is real. Furthermore, the behaviour of the AGN space density at high redshift will be determined when optical identification programs will be completed."
    },
    "0011/astro-ph0011177_arXiv.txt": {
        "abstract": "{  Recently much work in Gamma-Ray Burst (GRB) studies was devoted to revealing the nature of outburst mechanism and to looking for GRB progenitors.  Several types of progenitors were proposed for GRBs.  Most promising objects are collapsars, compact object binaries,  Helium star mergers and recently discussed supernovae. In this paper we consider four proposed binary star progenitors  of GRBs: double neutron star (NS-NS), black hole neutron star  (BH-NS), black hole white dwarf (BH-WD) mergers and Helium  star mergers (He-BH). Helium star mergers are a possible outcome of common envelope  evolution of a compact object entering the envelope of a giant  with a helium core.  Using population synthesis we calculate number of the binary progenitors and show that BH-WD and Helium star mergers dominate  population  of the proposed binary progenitors. Comparison of distribution of different binary mergers around galaxies they are born in, with localization of GRB afterglows in their host galaxies shows that only Helium mergers may be responsible for GRBs with observed afterglows while it excludes NS-NS, BH-NS and BH-WD systems as GRB progenitors. Assuming that all GRBs come from Helium star mergers and comparing  numbers of Helium star mergers with observed BATSE GRB rate let us  derive upper limit on GRB collimation to be $\\sim 4^\\circ$. } ",
        "introduction": "The last decade  brought a great breakthrough in gamma-ray burst  studies. Observations of GRB afterglows in X-ray, optical and radio wavelength  domains \\cite{1997IAUC.6576....1C,1997IAUC.6584....1G}  led to identification of GRB host galaxies \\cite{1997IAUC.6588....1G} and measurements of their redshifts. This has solved the long standing problem of their origin. While we learned that GRBs come from cosmological distances there  are still two major difficulties in understanding this phenomenon. First, we do not understand fully the physics of outburst. Although several models were proposed they all yet have to  meet some severe constraints imposed by observations (i.e. releasing  energies of $10^{51}$--$10^{54}$\\ ergs in timescales as short as $10^{-2}\\,$s in case of some GRBs). Second we don't know what are the astrophysical objects leading to  gamma-ray bursts, i.e. what are their progenitors?  Recently a black hole accretion disk model of GRB outburst has been  favored \\cite{Me99,Bl00,FWH99}. Progenitors leading to this model include  collapsars \\cite{Woosley93,MacFadyen99} and binary stars:  Helium star mergers \\cite{Fryer98},  double neutron stars \\cite{Ruffert97,Meszaros97},  black hole neutron star \\cite{Lee95,Kluzniak98} and  black hole white dwarf systems \\cite{Fryer99b}.  Also recently  the  supernovae  gamma-ray bursts connection received much attention \\cite{P99,W00,Ch00}, however there is  still no clear evidence that these two phenomena are intrinsically  correlated \\cite{1999A&AS..138..469G}.  The present paper is an extension of our previous studies \\cite{BB1998,BBZ99,BBZ2000} and we aim here at a discussion of viability of the proposed GRB binary progenitors.  One way of telling which group of proposed binaries might be responsible for GRBs is to predict their numbers and compare them to the observed number of GRBs.  Population synthesis is a powerful tool for predicting numbers of  binary populations although it suffers from many uncertainties as some parameters of single and binary evolution are poorly known.  Moreover, population synthesis works well in predicting  the relative numbers of events, while calculation of absolute numbers requires additional assumptions. Using population synthesis method we calculate production rates of Helium star mergers, double neutron stars, black hole neutron star and black hole white dwarf systems. We use BATSE detection limit of observed number of GRBs to compare  with our predicted numbers of the binary progenitors.  Another way of discerning among the binary progenitors is to compare their merger site distributions around host galaxies with localization of GRBs within host galaxies. Again this may be accomplished by population synthesis method to  simulate a given binary population and then placing it in  galactic gravitational  potentials one may trace each binary until its components merge due  to gravitational wave energy losses. We perform such  calculations for different galactic potentials and then compare the results to the observations of GRB afterglows and their  positions within host galaxies.  In this work we extend our previous studies to include two more proposed binary progenitors: black hole white dwarf binaries and  Helium star mergers. Moreover, we improve our population synthesis code and for consistency  we also present the updated results for two previously studied types of  proposed progenitors: double neutron star and black hole neutron  star systems.  Population synthesis method has already been applied to study compact  object binaries in context of GRB progenitors. However most of authors have concentrated only on double  neutron stars and black hole neutron star systems \\cite{Nar91,1991ApJ...380L..17P,1993MNRAS.260..675T,1996A&A...312..670P,1999MNRAS.305..763B,BB1998,BBZ99,BBZ2000}. Only one group \\cite{FWH99} presented calculations including all types  of proposed binary progenitors and also collapsars. Our calculations besides the conclusions we obtain, may serve as a direct check and  comparison for the  results presented by Fryer et al., \\cite*{FWH99}.  In section 2 we present observational data on the distribution of GRBs around their host galaxies, in section 3 we describe the population synthesis code used in this paper and formation scenario for Helium star mergers, in section  4 we discuss the results, and we present our conclusions in section 5. ",
        "conclusions": "We have calculated the properties of the possible binary GRB progenitors: BH-NS, NS-NS, BH-WD, and He mergers.  The  GRB binary progenitor production rates fall off exponentially with width of natal kick velocity distribution. We calculate the expected redshift distributions and numbers for each type of the progenitor, and find that BH-WD and He-BH type events dominate over the NS-NS or BH-NS mergers. Moreover,  in calculating the redshift distribution of binary progenitors of GRBs one can not neglect the evolutionary times $t_{evol}$. In our example calculation we find that  assuming that all GRBs result from binary mergers, then the population is dominated by BH-WD and He-BH star mergers, and the collimation must be of order {$\\sim 3 \\times 10^{-3}\\ (\\sim 4^\\circ)$}. The existence or non existence of GRBs at high redshifts depends on the star formation history. In particular the recent measurement of $z=4.50$ in case of GRB000131 \\cite{Andersen2000} suggests that the star formation rate has been high up to high $z$.  We find that only the  He-BH (Helium star merger) model of GRBs is consistent with Beppo SAX observations of long bursts and their localization in host galaxies. As noted in several previous papers a larger fraction of NS-NS binaries merge outside even massive host galaxies. The distribution of BH-NS mergers is tighter around the massive  galaxies, however we find that they do escape from the potentials of small galaxies. Our population synthesis code results indicate that  a significant fraction of the BH-WD binaries should obtain  large enough velocities so that they can escape from the potential well even of a massive galaxy. This differs  from the results of Fryer et al., \\cite*{FWH99}, who find that the  BH-WD binaries should lie within the host galaxies. Moreover our code results in a much larger fraction of  BH-WD binaries than the fraction Fryer et al., \\cite*{FWH99} obtain. Finally we remark that  NS-NS, BH-NS and BH-WD mergers can be responsible for short bursts which are not observed by Beppo SAX.  This could be resolved by  HETE-II, if short burst afterglows are discovered, and if their host galaxies are identified."
    },
    "0011/astro-ph0011207_arXiv.txt": {
        "abstract": "Unified recombination cross sections and rates are computed for (e~+~Fe~XVIII) $\\longrightarrow$ Fe~XVII including non-resonant and resonant (radiative and di-electronic recombination, RR and DR) processes in an ab initio manner with relativistic fine structure. The highly resolved theoretical cross sections exhibit considerably more resonance structures than observed in the heavy-ion storage ring measurements at Heidelberg, Germany. Nonetheless, the detailed resonance complexes agree well with experiment, and the unified rates agree with the sum of experimentally derived DR and theoretical RR rates to $\\sim$ 20\\%, within experimental or theoretical uncertainties. Theoretical results may provide estimates of field ionization of rydberg levels close to the DR peak, and non-resonant background contributions, particularly close to the RR peak as E $\\rightarrow$ 0. More generally, the unified results avoid the physical and practical problems in astrophysical models inherent in the separation of electronic recombination into RR and DR on the one hand, and further subdivision into low-energy $\\Delta n = 0$ DR and high-energy $\\Delta n > 0$ DR in photoionized and collisionally ionized X-ray plasmas. ",
        "introduction": "The {\\it Chandra X-ray Observatory} and the {\\it X-ray Multi-Mirror Mission - Newton} are observing a variety of sources (Canizares \\etal 2000, Predehl \\etal 2000) with wide-ranging plasma conditions. The analysis requires astrophysical models (e.g. Kallman 1995, Brickhouse \\etal 1995) whose accuracy depend on the atomic cross sections for collisional and radiative processes such as electron impact excitation, photoionization and recombination. K- and L-shell iron ions are among the most prominent atomic species in many sources such as active galactic nuclei (Ogle \\etal 2000), stellar coronae and winds (Schulz \\etal 2000), and cooling flows in clusters of galaxies (Fabian \\etal 2000). Excitation, photoionization, and recombination of Ne-like Fe~XVII are of particular interest as this ion is a strong X-ray radiator owing to a multitude of L-shell excitations at T $<$ 1 keV.  Although most of these atomic parameters are obtained theoretically, those need to be benchmarked against available experimental measurements. In recent years there has been considerable progress in both theoretical and experimental methods. High resolution electron-ion recombination measurements on ion storage rings show very detailed resonance structures in the low energy region usually accessible in experiments (e.g. Mannervik \\etal 1997, Kilgus \\etal 1992). Recently, such measurements have been done for (e~+~Fe~XVIII) $\\rightarrow$ Fe~XVII (Savin \\etal 1997,1999) in the low energy region dominated by $\\Delta$ n = 0 resonances. The experimental results naturally measured the combined non-resonant and resonant contributions. These were then processed to separate the background and extract the resonance contributions (DR) by fitting with an experimentally deduced beam shape function. Savin \\etal find that their inferred DR rates from low-energy measurements differ from previous theoretical calculations by up to a factor of 2 or more. Their best agreeement of $\\approx$ 30\\% is with multi-configuration Dirac-Fock (MCDF) and Breit-Pauli (MCBP) calculations in the isolated resonance approximation, assuming autoionization and the radiative probabilities independent of each other; detailed cross sections are not calculated. The comparison of derived DR rates from individual resonances (or blends) showed varying levels of agreement.  As the (e~+~ion) recombination  is unified in nature, it is theoretically desirable to consider the non-resonant and resonant processes (RR and DR) together. A unified theoretical formulation has been developed (e.g. Nahar and pradhan 1994, NP94; Zhang \\etal 1999, Z99), including relativistic fine structure (Zhang and Pradhan 1997), and used to compute cross sections and rates for many atomic systems, such as the K-shell systems C~IV - C~V and Fe~XXIV - Fe~XXV of interest in X-ray spectroscopy (Nahar \\etal 2000a,b; N99a,b). The unified results may be directly compared with experimental results, without the need to separate RR and DR. In this {\\it Letter}  we present new results for Fe~XVII and show that not only are the unified cross sections and rates in very good agreement with experiment, but may also be used to study important physical effects such as ionization of high-rydberg bound and autoionizing levels. More generally, the results demonstrate that the unified calculations avoid the basic inconsistency and incompleteness of photoionization and recombination data for modeling of laboratory and astrophysical plasma sources. ",
        "conclusions": "Astrophysical X-ray plasmas may be photoionized and/or collisionally ionized. Modeling of X-ray sources therefore requires total electronic recombination cross sections and rates at a wide range of energies and temperatures. To that end, we calibrate the unified theoretical and experimental data for recombination to Fe~XVII against each other; we find good agreement in the low-energy region accessible experimentally. All other previous unified BPRM calculations have also shown very good agreement with experiments for few-electron recombined ions C~IV, C~V, O~VII (Z99), Ar~XIV (Zhang and Pradhan 1997), and Fe~XXIV (Pradhan and Zhang 1997). Unified recombination cross sections including resonant and non-resonant processes (RR and DR) may therefore be computed for most astrophysially abundant elements.  In future, given the accuracy and resolution of the unified cross sections, and possible partial delineation according to n,$\\ell$, and J, an exact comparison with experimental measurements might yield precise information on: (i) the missing background non-resonant (RR) contribution near threshold as $E \\rightarrow 0$, into high rydberg levels, (ii) extraneous background contribution such as due to charge transfer, and (iii) high rydberg resonances contributing to the DR peak, and a more accurate field ionization cut-off than the approximate formulae heretofore employed.  Whereas the detailed comparison with experiment (Fig. 1) for low-energy $\\Delta n = 0$ (E $\\leq$ 140 eV) recombination is a useful test of accuracy, and relevant to low-T photoionized sources, collisionally ionized (coronal) high-T sources require total recombination rate coefficients up to T = $10^8$ K (Arnaud and Raymond 1982). That, in turn, requires recombination cross sections up to E $\\approx$ 1 keV or higher. In fact, we may predict that the next DR peak at about 800 eV due to the n = 3 levels will be much bigger than at $^2S_{1/2}$ (Fig. 1), since there are several strong dipole transitions with A-values about 1-2 orders of magnitude higher than $A(^2S_{1/2} \\rightarrow ^2P^o_{3/2,1/2})$. The first set of the these  resonances begins at just above 400 eV, close to the temperature of maximum abundance of Fe~XVII in coronal equilibrium. Correspondingly there will be a higher peak in the recombination rate than the one around 100 eV (Fig. 2). The $\\Delta n = 1$ DR therefore is expected to contribute more to total electronic recombination of (e~+~Fe~XVIII) $\\rightarrow$ Fe~XVII than the $\\Delta n = 0$ core transitions. The higher energy unified recombination calculations are in progress.  One of the main points is that since photoionization and recombination are treated as inverse processes with the same eigenfuncation expansion for the core ion (Eq. 1), the same set of resonances and non-resonant background are included in both processes - an essential requirement for self-consistency in photoionization equilibrium, as expressed by, \\begin{equation} \\int_{\\nu_0}^{\\infty} \\frac{4 \\pi J_{\\nu}}{h\\nu} N(X^{z}) \\sigma_{PI}(\\nu,X^{z}) d\\nu = \\sum_j N_e N(X^{z+1}) \\alpha_R(X_j^{z};T), \\end{equation} where $\\sigma_{PI}$ is the photoionization cross section and $J_{\\nu}$ the radiation flux. In general, the sum on the RHS of Eq. (2) extends over the infinite number of recombined bound levels. Now if the $\\alpha_R$ on the RHS is subdivided into non-resonant RR and resonant DR, as in existing photoionization models, then fundamental inconsistencies result. The RR rate is suppoesedly derived from non-resonant photoionization cross sections. But as we see from Fig. 1, the near threshold region is dominated by resonances and a rising non-resonant background. Therefore, photoionization calculations must also be carried out including the same resonances (c.f. the Opacity Project work; Seaton \\etal 1994, and references therein). There is however the issue of accurate resonance positions, which may be ameliorated by (a) averaging over the radiation field on the LHS of Eq. (2), and (b) pre-averaging over resonances in photoionization cross sections (Bautista \\etal 1998). The  further sub-division of DR into $\\Delta n = 0$ rates (as, for example, derived experimentally by Savin \\etal 1999) appropriate for only for low-T plasmas, and $\\Delta n > 0$ DR needed in high-T plasmas, implies that the recombination rates must be obtained for (RR + DR ($\\Delta n = 0$) + DR ($\\Delta n > 0$)), generally using different approximations and possibly valid in different temperature regimes. These problems: (i) inconsistent photoionization and recombination, (ii)unphysical division of RR and DR, and (iii) low and high energy DR in different but overlapping energy(temperature) ranges, may be overcome with the unified method for electronic recombination, and corresponding photoionization cross sections, to enable a self-consistent treatment of photoionization and recombination in X-ray photoionized sources."
    },
    "0011/astro-ph0011031_arXiv.txt": {
        "abstract": "I review simultaneous constraints on the cosmological parameters $\\lambda_{0}$ and $\\Omega_{0}$ from gravitational lensing.  The emphasis is on systematic extragalactic surveys for strong gravitational lenses, mainly the largest and best-defined such survey, JVAS/CLASS. ",
        "introduction": "Since the details of the gravitational-lensing effect depend on the cosmological model, it offers a means of determining the cosmological parameters $H_{0}$, $\\lambda_{0}$ and $\\Omega_{0}$ by comparing expectations from different cosmological models with observations. Advantages of using gravitational lensing to learn about the cosmological model include the fact that it is based on relatively well understood astrophysics and that it makes use of information from an intermediate redshift range, complementing tests which use information primarily from the low-redshift (e.g.\\ cosmic flows) or high-redshift (e.g.\\ cosmic microwave background fluctuations) regimes.  I review simultaneous constraints on the cosmological parameters $\\lambda_{0}$ and $\\Omega_{0}$ from gravitational lensing.  (Constraints on $H_{0}$ are discussed in other contributions in this volume; in particular, Schechter discusses constraints from lensing.)  The emphasis is on systematic extragalactic surveys for strong gravitational lenses, mainly the largest and best-defined such survey, JVAS/CLASS. However, other methods of constraining $\\lambda_{0}$ using (mostly strong) gravitational lensing are also discussed.  After briefly reviewing the basic theory and history of the subject, I present the currently available constraints and briefly (since this is covered in other contributions in this volume, such as that by Lineweaver) touch on joint constraints with other cosmological tests. Finally, I discuss systematic errors and prospects for the future.  In general, one can derive constraints on $\\lambda_{0}$ and $\\Omega_{0}$ from gravitational lensing when there is more than one source plane involved (e.g. Golse, Kneib, \\& Soucail 2000).  This is the case in most examples of weak lensing and cluster lensing.  Also, arc statistics (e.g.\\ Bartelmann et al. 1998) can give information about the cosmological model, not only since more than one source plane is involved, but also since this is sensitive to the evolution of structure (the lenses in this context), which also depends on the cosmological model.  Another possibility is provided by higher-order effects in measuring $H_{0}$ through gravitational-lens time delays; the time delay is inversely proportional to $H_{0}$, but there is a weaker, non-linear dependence on $\\lambda_{0}$ and $\\Omega_{0}$.  Finally, the statistical analysis of gravitational-lens surveys provides a potentially powerful method of measuring the cosmological parameters. ",
        "conclusions": "\\begin{itemize} \\item Despite vast improvements on both the observational and theoretical sides, the upper limit on $\\lambda_{0}$ has been a remarkably stable number in the astronomical literature. \\item We can already place a {\\em statistically\\/} strong upper limit on $\\lambda_{0}$. \\item Gravitational-lensing statistics is already hinting at a positive lower limit on $\\lambda_{0}$, at least in a flat universe.  While there is abundant evidence for a positive cosmological constant based on various combinations of cosmological tests, at present only the $m$-$z$ relation for type Ia supernovae (see e.g. the contributions by Perlmutter and Kirschner in these proceedings) is the only cosmological test which indicates this {\\em by itself}; it would thus be nice to have another test which can do so. \\item When the analysis of the CLASS gravitational-lens survey is complete, this will improve in the sense that the statistical errors will become smaller and systematic effects will be reduced; whether or not the upper limit on $\\lambda_{0}$ decreases depends, of course, on what the value of $\\lambda_{0}$ actually is and to what extent current estimates are biased. \\item We will have a {\\em robust\\/} upper limit on $\\lambda_{0}$\\ soon when the $S$-$z$ plane for CLASS is better understood.  The $S$-$z$ plane enters into the calculation in two respects, since it determines both the flux-density dependent redshift distribution, which is needed to estimate the redshifts of the non-lensed sources in CLASS, and the redshift-dependent luminosity function, needed for the calculation of the amplification bias.  (See McKean et al., these proceedings, for some information about current work in this area.) \\item We will be able to provide an {\\em independent\\/} check on constraints from $m$-$z$ relation (SN Ia), CMB, (evolution of) LSS and clusters, weak lensing, cluster lensing etc. \\item Looking farther ahead, many more lens systems will be found in the future in both radio and optical surveys.  Since the Poisson noise due to the small number of systems is at present a source of appreciable uncertainty, progress can be expected just from finding more lens systems (after doing the corresponding analysis, of course). \\item Even when the cosmological parameters are known, perhaps more precisely through (a combination of) other methods, gravitational-lensing statistics will still be interesting:  It will be possible to use gravitational-lensing statistics to constrain variables other than $\\lambda_{0}$ and $\\Omega_{0}$, and thus study galaxy evolution, the source population etc. \\end{itemize}"
    },
    "0011/astro-ph0011082_arXiv.txt": {
        "abstract": "We present $RIJHK$ imaging of seven radio galaxies from the 7C Redshift Survey (7CRS) which lack strong emission lines and we use these data to investigate their spectral energy distributions (SEDs) with models which constrain their redshifts. Six of these seven galaxies have extremely red colours ($R-K>5.5$) and we find that almost all of them lie in the redshift range $1< z < 2$. We also present near-infrared spectroscopy of these galaxies which demonstrate that their SEDs are not dominated by emission lines, although tentative lines, consistent with H$\\alpha$ at $z=1.45$ and $z=1.61$, are found in two objects. Although the red colours of the 7CRS galaxies can formally be explained by stellar populations which are either very old or young and heavily reddened, independent evidence favours the former hypothesis.  At $z \\sim 1.5$ at least $1/4$ of {\\em powerful} radio jets are triggered in massive ($> L^{\\star}$) galaxies which formed the bulk of their stars several Gyr earlier, that is at epochs corresponding to redshifts $z \\gtsimeq 5$. If a similar fraction of all $z \\sim 1.5$ radio galaxies are old, then extrapolation of the radio luminosity function shows that, depending on the radio source lifetimes, between 10 -- 100 \\% of the near-IR selected extremely red object (ERO) population undergo a radio outburst at epochs corresponding to $1 < z < 2$. An ERO found serendipitously in the field of one of the 7CRS radio sources appears to be a radio-quiet analogue of the 7CRS EROs with an emission line likely to be [OII] at $z=1.20$.  The implication is that some of the most massive elliptical galaxies formed the bulk of their stars at $z \\gtsimeq 5$ and these objects probably undergo at least two periods of AGN activity: one at high redshift during which the black hole forms and another one at an epoch corresponding to $z \\sim 1.5$ . ",
        "introduction": "At both low and high redshifts, powerful radio sources are believed to reside exclusively in giant elliptical galaxies or their progenitors. Surveys of radio galaxies and radio-loud quasars can therefore be used to trace the evolution of such massive galaxies over cosmological timescales. 3C radio galaxies at $z \\gtsimeq 0.6$ show extended continuum and emission line structures aligned along their radio axes (McCarthy et al. 1987; Chambers et al. 1987). Investigation of the cause of this so-called `alignment effect' has shown that, while this is not due to the bulk of the stellar population, a single mechanism, e.g. recent star-formation or dust-scattered quasar light, cannot account for all cases (Best et al. 1997). Best et al. and Blundell \\& Rawlings (1999) have pointed out the importance of factoring in source age to the interpretation of this. However, it is clear that the strength of the alignment effect decreases with decreasing radio luminosity (Lacy et al. 1999a) and is weaker in the observed-frame near-infrared than in the optical (Dunlop \\& Peacock 1993; Eales et al. 1997). Dunlop \\& Peacock (1993) found that the $R-K$ colours of lower power radio galaxies are consistent with evolved stellar populations. The same is true of the host galaxies of 3C sources at $z\\sim 1$, once the aligned component has been taken into account (Best et al. 1998).  The reddest $z \\sim 1$ radio galaxy in the 3CRR sample of Laing, Riley \\& Longair (1983) is 3C 65. The light from this object has been shown to be dominated by a very old stellar population ($\\sim 4$ Gyr; Lacy et al. 1995; Stockton, Kellogg \\& Ridgway 1995) although a reddened quasar nucleus seems to make some contribution to its red colour (Lacy et al. 1995; Simpson, Rawlings \\& Lacy 1999).  Two extremely red radio galaxies at $z \\sim 1.5$ have subsequently been discovered in the follow-up of the faint Leiden Berkeley Deep Survey (LBDS 53W091 --- Dunlop et al. 1996; LBDS 53W069 --- Dunlop 1999).  Keck spectroscopy of these galaxies shows that their optical continua are dominated by very old stellar populations: ages $\\gtsimeq 3$ Gyr according to Spinrad et al. (1997) and Dunlop (1999) although ages $\\sim 1.5 ~ \\rm Gyr$ are preferred by Yi et al. (2000). These observations indicate very high formation redshifts ($z \\gtsimeq 5$) for the bulk of stars in such galaxies. Peacock et al. (1998) showed that such high formation redshifts are consistent with our current knowledge of the cosmic power spectrum on $\\sim$ Mpc scales.  Near-infrared (near-IR) surveys have discovered a population of galaxies which are very faint or undetected at optical wavelengths (e.g. Elston, Reike \\& Reike 1988, 1989; Hu \\& Ridgway 1994). Galaxies with observed colours of $R-K \\geq 6$ have come to be known as {\\em Extremely Red Objects} (EROs). With the advent of wide-field near-IR detectors it is now possible to detect large numbers of such objects. There are generally two ways in which such red optical to near-IR colours can be achieved; an evolved stellar population at $z > 1$ or a highly reddened, dusty starburst or AGN. Much activity is currently focussed on trying to understand which of these two types of object dominate the observed EROs. As yet, there are only a few clear determinations for individual sources. The ERO HR 10 has been found to be a very dusty, ultra-luminous infrared galaxy at $z=1.44$ (Dey et al. 1999). At the time of writing this is the only near-IR selected ERO which has been shown to have its extreme colours due to dust, however some of the optical counterparts of sub-mm blank-field surveys are also found to be extremely red (Smail et al. 1999). In contrast, there are a couple of EROs which appear to be unreddened evolved stellar populations at $z \\approx 1.5$ (CL0939+4713B - Soifer et al. 1999; HDFS 223251-603910 - Stiavelli et al. 1999). Near-IR spectroscopy by Cimatti et al. (1999) has shown the majority of a sample of nine red galaxies with $R-K>5$ appear to be old stellar populations at $0.8 < z < 1.6$, with only a couple of cases of possible dusty starbursts. Studies with the Hubble Space Telescope (HST) show the majority of EROs to have E/S0 morphologies (Moriondo, Cimatti \\& Daddi 2000; Stiavelli \\& Treu 200). The detection of a strong clustering signal in the ERO population (Daddi et al. 2000a) is in agreement with the hypothesis that massive evolved galaxies dominate the population.  The exact nature of extremely red galaxies is an important test of models of galaxy formation and evolution. A high space density of EROs which are evolved, elliptical galaxies at intermediate redshifts would mean high formation redshifts for most massive galaxies. However, if the EROs are due to ongoing dusty star-formation, possibly in mergers, it would support certain hierarchical models for galaxy formation. The likely redshift range of EROs in near-IR samples limited to $K (2.2 \\mu{\\rm m}) \\sim 19$ is $1 \\ltsimeq z \\ltsimeq 2$. The main reason for the lower redshift limit is that the continuum below the 4000 \\AA~ break passes into the observed $R$-band, while the upper redshift limit is the result of an upper limit to the intrinsic stellar luminosity of a galaxy. Note that the radio-selected EROs from the LBDS survey have a co-moving space density $\\sim 100$ times lower than that of the near-IR selected EROs.  We have been working on securing optical identifications and redshifts for a complete sample of 76 radio sources selected at a low radio frequency, namely 151 MHz.  The flux-density limit for this sample --- the combination of the 7C-I and 7C-II regions of the 7C Redshift Survey (hereafter referred to as the 7CRS) --- is $S_{151} \\geq 0.5$ Jy and the total sky area covered by regions I and II is 0.013 sr. The flux-density limit is a factor of 25 lower than the 3CRR sample and a factor of 40 higher than the 2 mJy 1.4 GHz LBDS survey (assuming a radio spectral index of 0.8). The lower flux-density selection than the 3CRR sample means we expect little contribution from non-stellar emission to the optical continua of 7CRS radio galaxies (c.f. Lacy et al. 1999a). All the sources have been identified at $K$-band and optical spectroscopy has led to redshifts for 90\\% of the sample (Willott et al., in prep. and Blundell et al., in prep., will present full details of the 7C-I and 7C-II samples).  The seven sources for which optical spectroscopy did not secure a redshift are very faint or invisible in our optical images and have red optical-to-near-IR colours. Four of them are bona-fide EROs with $R-K>6$, of which the most extreme example is undetected at $R$-band with $R-K>7.3$. Since complete redshift information is crucial for deducing the nature and evolution of radio sources, we have obtained multi-colour photometry for these objects in order to model their spectral energy distributions (SEDs) and hence constrain their redshifts. In addition, we have obtained some low-resolution near-infrared spectroscopy to search for emission lines and/or breaks in the continuum.  In this paper, we use these observations to constrain the redshifts of these sources and attempt to determine whether their extremely red colours are due to old stellar populations or to the effects of dust reddening on a younger population. In Section 2 we list the objects studied in this paper and discuss their colours in the context of the sources with spectroscopic redshifts in the 7CRS. In Section 3 we present our observations and in Section 4 we describe the SED fitting technique used. A discussion of the SED fitting for individual objects follows in Section 5. In Section 6 we review the reliability of the redshift estimation. In Section 7 we discuss the implications of our findings with respect to radio galaxies, the field ERO population and the galaxy population in general. We report on the serendipitous discovery of a radio-quiet ERO with a spectroscopic redshift of $z=1.200$ in the appendix. We assume that $H_0=50~ {\\rm km~s^{-1}Mpc^{-1}}$ and consider two spatially-flat cosmologies: $\\Omega_{\\mathrm M} =1$, $\\Omega_ \\Lambda=0$ and $\\Omega_{\\mathrm M}=0.3$, $\\Omega_\\Lambda=0.7$.  \\begin{figure*} \\epsfxsize=0.95\\textwidth \\epsfbox{fig1.ps} {\\caption[junk]{\\label{fig:rkz} Observed $R-K$ colour as a function of redshift for sources in the 7C-I and 7C-II regions of the 7C Redshift Survey (7CRS). Filled circles are quasars, crossed circles are broad-lined radio galaxies, open circles are high emission line luminosity radio galaxies and open triangles are low emission line luminosity radio galaxies (see Willott et al. 1998, 2000a for definitions of these categories). The open squares are the seven 7CRS objects without redshifts from optical spectroscopy which have redshifts estimated in this paper. The asterisks show the two red radio galaxies from the LBDS survey which have redshifts inferred from stellar absorption lines (Dunlop 1999).  The solid lines show the evolution of the expected observed colour of a galaxy which formed in an instantaneous starburst at redshift $z=10$ for $\\Omega_{\\mathrm M} =1$, $\\Omega_ \\Lambda=0$ and for $\\Omega_{\\mathrm M}=0.3$, $\\Omega_\\Lambda=0.7$ using the models of Bruzual \\& Charlot (in prep.). The dotted line shows the change in $R-K$ as a function of $z$ for a typical unreddened quasar spectrum.  }} \\end{figure*} ",
        "conclusions": "\\begin{itemize}  \\item{We have shown that the optical to near-infrared SEDs of optically faint 7CRS radio sources are well fit by old or dusty galaxies in the redshift range $1 \\leq z \\leq 2$.}  \\item{Two galaxies show narrow emission lines in their near-IR spectra which if identified as H$\\alpha$ give redshifts of $z \\sim 1.5$ compatible with the SED fitting. The fluxes of these lines and the lack of lines in other galaxies are as expected from the correlation between narrow line and radio luminosities for radio sources.}  \\item{The reason most of these objects did not have redshifts from optical spectroscopy is that they fall within the redshift range where there are no bright lines in the optical wavelength region}.  \\item{Independent evidence suggests that the red colours of these galaxies are not caused by dust-reddening, so they are most likely very old elliptical galaxies which formed at very early epochs. These objects account for $\\approx 1/4$ of the powerful radio source population at these redshifts. It is possible that the other radio sources are blue due to non-stellar emission and have underlying evolved host galaxies, in which case the space density of radio galaxies with old stellar populations is a factor of 4 greater than determined here.}  \\item{We find that $\\sim 3\\%$ of near-IR selected EROs are expected to be radio-loud AGN with radio luminosities $L_{151} \\gtsimeq 10^{24} ~{\\rm W Hz}^{-1} {\\rm sr}^{-1}$. Depending upon the duration of radio activity it is possible that a large fraction ($10-100$\\%) of near-IR selected EROs go through such a period of radio activity at some point in their intermediate-redshift ($1 < z <2$) lifetimes.}  \\item{We have discovered a radio-quiet extremely red galaxy ($R-K=5.7$) ERO J081539+2446.8 in the field of 5C7.47 with an emission line in its optical spectrum. This object is best explained as a radio-quiet analogue of the 7CRS EROs at $z=1.200$ (see appendix). }  \\item{A consistent picture emerges that some of the most massive galaxies formed in violent starbursts with associated AGN activity at high redshift ($z > 2$) and then underwent a second bout of AGN activity at intermediate redshifts ($1<z<2$) but with only a small amount of further star-formation.}  \\end{itemize}"
    },
    "0011/astro-ph0011427_arXiv.txt": {
        "abstract": "Recent observations with the {\\it Submillimeter Wave Astronomy Satellite} indicate abundances of gaseous H$_2$O and O$_2$ in dense molecular clouds which are significantly lower than found in standard homogeneous chemistry models.  We present here results for the thermal and chemical balance of inhomogeneous molecular clouds exposed to ultraviolet radiation in which the abundances of H$_2$O and O$_2$ are computed for various density distributions, radiation field strengths and geometries.  It is found that an inhomogeneous density distribution lowers the column densities of H$_2$O and O$_2$ compared to the homogeneous case by more than an order of magnitude at the same $A_{\\rm V}$. O$_2$ is particularly sensitive to the penetrating ultraviolet radiation, more so than H$_2$O. The S~140 and $\\rho$ Oph clouds are studied as relevant test cases of star-forming and quiescent regions. The SWAS results of S~140 can be accommodated naturally in a clumpy model with mean density of $2\\times 10^3$ cm$^{-3}$ and enhancement $I_{\\rm UV}=140$ compared with the average interstellar radiation field, in agreement  with observations of [C I] and $^{(13)}$CO of this cloud.  Additional radiative transfer computations suggest that this diffuse H$_2$O component is warm, $\\sim 60-90$ K, and can account for the bulk of the $1_{10}-1_{01}$ line emission observed by SWAS. The $\\rho$~Oph model yields consistent O$_2$ abundances but too much H$_2$O, even for [C]/[O]=0.94, if $I_{\\rm UV}<10$ respectively $<40$ for a mean density of $10^3$ respectively $10^4$ cm$^{-3}$. It is concluded that enhanced photodissociation in clumpy regions can explain the low H$_2$O and O$_2$ abundances and emissivities found in the large SWAS beam for extended molecular clouds, but that additional freeze-out of oxygen onto grains is needed in dense cold cores. ",
        "introduction": "In contrast with the major carbon-bearing species (CO, C, C$^+$), the abundances of the major oxygen-containing atoms and molecules in general molecular clouds are still poorly understood. The best-known species is CO, which reaches large abundances of $\\sim 10^{-4}$ when the shielding column of gas exceeds $A_V\\approx$1 for a typical interstellar cloud of $\\sim 10^3$ cm$^{-3}$ illuminated by the average interstellar radiation field, locking up the bulk of the gas-phase carbon (e.g., Hollenbach \\& Tielens 1999, van Dishoeck \\& Black 1988).  Theoretical models indicate that at the edge atomic O is abundant, but that further into the cloud, more and more of the ultraviolet radiation is attenuated so that at $A_V\\approx 4$ mag, CO ceases to be the dominant oxygen-bearing species and O$_2$ reaches its peak abundance of $>10^{-4}$, with H$_2$O present at the level of $\\sim 10^{-6}$ (e.g., Tielens \\& Hollenbach 1985, Lee et al. 1996). If true, then H$_2$O and O$_2$ could be major coolants in molecular clouds (Goldsmith \\& Langer 1978; Neufeld, Lepp \\& Melnick 1995).   The launch of the {\\it Submillimeter Wave Astronomy Satellite} (SWAS) (Melnick et al.\\ 2000) has made it possible to look for emission lines of H$_2$O and O$_2$ in cold molecular clouds. Although the {\\it Infrared Space Observatory} (ISO) observed many H$_2$O lines, it was sensitive primarily to warm, $\\sim 100$ K molecular gas near young stellar objects where the H$_2$O abundance can be enhanced due to ice evaporation and high-temperature (shock) reactions (e.g., van Dishoeck et al.\\ 1999, Nisini et al.\\ 2000).  Measurements of, or upper limits on, the abundances of H$_2$O and O$_2$ provide crucial tests of many aspects of the oxygen chemistry in normal molecular clouds and of the cloud structure, and hence of our theoretical understanding of such systems.  The initial SWAS results indicate that the abundances of H$_2$O and O$_2$ in dense molecular cloud cores are surprisingly low compared with the above model predictions, about $6\\times 10^{-10} - 1\\times 10^{-8}$ for H$_2$O and $<$ a few $\\times 10^{-7}$ for O$_2$ (Ashby et al.\\ 2000, Snell et al.\\ 2000a, Goldsmith et al.\\ 2000). In contrast, ISO observations of the [O I] 63 and 145 $\\mu$m lines suggest substantial abundances of atomic O in quiescent gas (e.g., Baluteau et al.\\ 1997, Caux et al.\\ 1999).  Bergin et al.\\ (2000) propose a model in which freeze-out of oxygen on dust grains in the form of molecular ices is significant, with the remaining oxygen in atomic form in the gas, along with CO. While freeze-out and gas-grain interactions undoubtedly play a role in dense cold clouds, we present here an alternative explanation based on clumpy molecular clouds which may be more appropriate for the general lower density molecular cloud material contained in the large SWAS beam of $3.3'\\times 4.5'$. The inhomogeneous nature of molecular clouds is well established from observations of extended [C I] and [C~II] (e.g., Keene et al.\\ 1995, Stutzki et al.\\ 1988, Plume et al.\\ 1999), and has been modeled by various groups (e.g., Meixner \\& Tielens 1993, Spaans 1996, St\\\"orzer et al.\\ 1996). It is therefore natural to also explore the H$_2$O and O$_2$ abundances and emission in such models. ",
        "conclusions": "\\subsection{H$_2$O and O$_2$ abundances}  The observational results represent abundances integrated over depth. Thus, mass-weighted abundances for the spherical models have been computed for comparison with observations. Specifically, the mean abundance $X(r)$ enclosed within some fractional radius $r$ is given by integration of the abundance $A(r)$ through $X(r)$=$\\int_0^r A(r')r'^2dr'$/$\\int_0^r r'^2dr'$, where 0 is at the center and 1 at the edge of the model cloud. Note that only $1/8$ of the mass is contained inside the half radius of the spherical constant density model cloud, while this would be $1/2$ for a $1/r^2$ density profile.  For the clumpy, spherical model of Figure 1 (top and middle) the above approach leads to a predicted mean H$_2$O abundance $X(1)=3\\times 10^{-8}$ and an O$_2$ abundance of $4\\times 10^{-10}$ for a typical total depth of $A_V=10$~mag. Snell et al.\\ (2000a) detected H$_2$O emission from the S~140 dense star-forming core and PDR region with SWAS, but not from the surrounding extended molecular cloud.  Their inferred abundances of $\\sim 10^{-8}$ (see also the detailed analysis by Ashby et al.\\ 2000) for the core and $< 10^{-8}$ for the extended cloud are only a factor of three lower than our model results and consistent within the factor of three or more uncertainty in their assumed density (the inferred H$_2$O abundance scales with the inverse of density). Figure 1c (bottom) shows the effect of increasing the density to $5\\times 10^4$ cm$^{-3}$ for the S~140 environment.  The smaller value of $U$ causes photodissociation to play a lesser role in the removal of O$_2$ and H$_2$O, and hence clumpiness is of less importance. The model (mass weighted) abundances are $3\\times 10^{-7}$ and $4\\times 10^{-8}$ for H$_2$O and O$_2$, respectively.   The model for $\\rho$ Oph appears less successful in reproducing the observed abundances of O$_2$ and H$_2$O. The mass weighted H$_2$O abundance for Figure 2c (bottom) is $3\\times 10^{-7}$, a factor of $\\sim$100 above the measured value.  Panels 2a and b (top and middle), which are representative of larger values of $U$ and roughly consistent with the results of Liseau et al.\\ (1999) who found $n_{\\rm H}=3\\times 10^3-3\\times 10^4$ cm$^{-3}$, yield mass-weighted H$_2$O abundances of $5\\times 10^{-8}$ and $2\\times 10^{-8}$, respectively. From the discussion in \\S 3, it follows that the S~140 case yields better agreement with the observations mainly because of its larger illumination (and $U$ parameter). The clumpy models enhance the importance of photodissociation of H$_2$O and O$_2$ deeper into the cloud, and therefore work best for a strong incident radiation field.   In all inhomogeneous models with $n\\approx 10^3$ cm$^{-3}$, the mass weighted abundance of O$_2$ is below $10^{-7}$, both for S~140 and $\\rho$ Oph, consistent with the overall SWAS upper limits as well as the specific upper bounds for S~140, $<7\\times 10^{-7}$, and $\\rho$ Oph, $<3\\times 10^{-7}$ (Goldsmith et al.\\ 2000). Only the quiescent model of Figure 2c (bottom) with high density and low $I_{\\rm UV}$, yields a mass weighted O$_2$ abundance of $10^{-6}$. This is a factor of a few above the observational limit (Goldsmith et al.\\ 2000), but this is easily alleviated by a modest increase in the [C]/[O] ratio (Bergin et al.\\ 2000). To test the latter possibility, we have run a model with an oxygen depletion of 0.17 and [C]/[O]=0.94. A decrease in the mass weighted H$_2$O (O$_2$) abundance of a factor of 5 (20) is found, bringing the model in agreement with the observations for S~140 and factors of 20, 4 and 2 too high for H$_2$O toward $\\rho$ Oph in models 2a, b and c, respectively, although a smaller value for the H$_3$O$^+$ branching ratio of 0.05 instead of 0.3 could further reduce these discrepancies.  \\subsection{H$_2$O and O$_2$ emission}  In order to verify whether the low abundances of H$_2$O and O$_2$ found for the clumpy model clouds can also reproduce the bulk of the observed  emission, radiative transfer computations have been performed using a Monte Carlo method (Spaans 1996). It was found that two effects dominate the emissivity of the H$_2$O $1_{10}-1_{01}$  line: temperature and geometry. The temperature is very important because it influences both the ortho-para ratio of the main collision partner H$_2$, 0.1 at 30 K and 1.0 at 80 K (Neufeld \\& Sternberg 1999), as well as the collisional excitation rate. The H$_2$ ortho-para ratio is a crucial parameter to compute rigorously from the thermal balance and FUV pumping since the o-H$_2$ $J$=1 collision rates with H$_2$O are more than an order of magnitude larger than the p-H$_2$ $J$=0 rates (Phillips, Maluendes \\& Green 1996). The geometry is important because the $1_{10}-1_{01}$ line is optically thick but effectively thin. This implies that line photons emitted by the warm edges of a spherical clump can interact with water molecules in the colder interior, and so constitute an excitation term that raises the level population in the $1_{10}$ state. Note in this that the curvature of the sphere causes the warm gas to cover $4\\pi$ steridians, unlike the slab case. It is found that a spherical geometry, for the S~140 temperature profile given below, raises the $1_{10}-1_{01}$ line strength with a factor of 2.4 compared to that from the plane-parallel slab. Furthermore, for identical physical conditions, a plane-parallel slab model yields a temperature profile that is systematically 20-30\\% cooler than for the corresponding spherical case, mainly because the latter possesses more lower extinction lines of sight.  The thermal balance yields a temperature of 150 (80) K at the edge of S~140 ($\\rho$~Oph) and 20 (10) K at its center. The mass weighted H$_2$O temperatures are $\\sim 80$ K and $\\sim 30$ K for S~140 and $\\rho$~Oph, respectively. For S~140, these values refer to the model of Fig.\\ $1a/b$, and for $\\rho$~Oph to the model of Fig.\\ $2b$. The S~140 value is significantly larger than the typical temperature in dense cores, $\\sim$ 30~K. The resulting integrated line intensities convolved with the SWAS beam are 1.6 K km s$^{-1}$ and 3.8 K km s$^{-1}$ for S~140 and $\\rho$~Oph, respectively. The clumpy S~140 model can thus reproduce the observed value (1.7 K km s$^{-1}$) quite nicely, but the clumpy $\\rho$~Oph H$_2$O emission, just as the abundance, is too large compared to the observed value of 0.8 K km s$^{-1}$. Use of a plane-parallel geometry would reduce the discrepancy for $\\rho$ Oph significantly. In both instances, the bulk of the signal in the SWAS beam is associated with the PDR. In fact, {\\it the H$_2$O emission arises mainly from the warm, $\\sim$60-90 K, edges of irradiated clumps}. Therefore, a strong correlation should exist between the SWAS H$_2$O emission and high$-J$ (e.g.\\ 5-4, 6-5) CO emission. Such a correlation is in fact observed toward M~17SW (Snell et al.\\ 2000b). It is the elevated temperatures at the surfaces of the clumps that makes the results so sensitive to the ortho-para ratio. For the case of S~140, the extended molecular cloud gives an emission that is a factor of 30 weaker at a position $4'$ north-east of the PDR, due to the lower ambient temperature, consistent with observations. The O$_2$ $N_J=3_3-1_2$ line is truly optically thin and yields, due to its simpler excitation, more robust results. We find $1.8\\times 10^{-5}$ K km s$^{-1}$ and $2.3\\times 10^{-3}$ K km s$^{-1}$ for the S~140 and $\\rho$~Oph case, respectively.   ISO H$_2$O observations suggest that much of H$_2$O toward star-forming cores such as S~140 is associated with a small region with enhanced H$_2$O surrounding the young stellar object(s) (e.g., Wright et al.\\ 1997, Ceccarelli et al.\\ 1999). Since this emission is subsequently beam-diluted in the large SWAS beam, its contribution could be small and still be consistent with the fact that the lower density gas in the SWAS beam is capable of reproducing the bulk of the detected water emission by SWAS. We have verified this by adding a clumpy star-forming core with $n=10^4$ cm$^{-3}$ and a size of $\\sim 0.5$ pc (Zhou et al.\\ 1994), and the same values for $f$ and $r$, to the S~140 extended molecular cloud as in Spaans \\& van Dishoeck (1997). Even though the mean water abundance goes up by a factor of 30, the lower kinetic temperature and smaller spatial extent yield a relative contribution of $\\sim 20$\\% after convolution with the SWAS beam. Furthermore, the Zhou et al.\\ CS observations can be well reproduced by this clumpy core that includes densities of up to $\\sim 10^5$ cm$^{-3}$ and temperatures that are still $\\sim 30-40$ K. \\footnote{At even higher densities ($\\sim 10^7$ cm$^{-3}$) and smaller scales ($\\sim 200-400$ AU), the H$_2$O abundance becomes even larger and the region emits thermally at about 100 K. Strong beam-dilution limits this contribution to the SWAS signal to less than 2\\%  at the distance of $\\sim 910$ pc for S~140.}     Bergin et al.\\ (2000) have argued against the importance of photodissociation because it would affect the abundances of other species such as NH$_3$ as well.  The S~140 model, including the star-forming core as above, was checked for consistency with the observational NH$_3$ values of $2-9\\times 10^{14}$ cm$^{-2}$ (Ungerechts, Winnewisser \\& Walmsley 1986). Our model value of $\\sim 4\\times 10^{14}$ cm$^{-2}$ is not in conflict with observations. In addition, even a moderate enhancement in NH$_3$ formation due to grain-surface chemistry is likely to resolve any mild discrepancy.  For the extended molecular cloud outside the core there are no observational constraints for ammonia. Future observational and theoretical studies should investigate other species whose abundances are sensitive to photodissociation, such as the CN/HCN and C$_2$H/HCO$^+$ ratios (e.g., Jansen et al.\\ 1995). This work has not addressed the effects of time-dependent chemistry and freeze-out on the O$_2$ and H$_2$O abundances (Bergin et al.\\ 2000, Viti 2000, private communication). Our results suggest that clumpiness is a viable alternative explanation for extended lower-density molecular clouds in star-forming regions like S~140 with high radiation field, but they also strengthen the case for time-dependent chemistry and freeze-out for a more quiescent dense region like $\\rho$ Oph. The freeze-out models may in fact benefit from adoption of a clumpy structure since the central dust temperature is enhanced by the higher internal ultraviolet field, thus providing a mechanism to moderate the very rapid freeze-out. Future observatories such as ODIN and the HIFI instrument on the {\\it Far Infrared and Submillimeter Telescope} (FIRST) will have increased sensitivity, and, in the case of FIRST, a much smaller beam. Therefore, more definite tests of the various classes of models will be forthcoming in this decennium."
    },
    "0011/astro-ph0011382.txt": {
        "abstract": "We propose that the curious Galactic Center filament known as ``The Snake'' is a twisted giant magnetic flux tube, anchored in rotating molecular clouds. The MHD kink instability generates coils in the tube and subsequent magnetic reconnection  injects relativistic electrons. Electrons diffuse away from a coil at an energy-dependent rate producing a flat spectral index at large distances from it. Our fit to the data of \\citet{gray95a} shows that the magnetic field $\\sim 0.4 \\> \\rm mG$ is large compared to the ambient $\\sim 7 \\mu \\> \\rm G$ field, indicating that the flux tube is force-free. If the {\\em relative} level of turbulence in the Snake and the general interstellar medium are similar, then electrons have been diffusing in the Snake for about $3 \\times 10^5 \\> \\rm yr$, comparable to the timescale at which magnetic energy is annihilated in the major kink. Estimates of the magnetic field in the G359.19-0.05 molecular complex are similar to our estimate of the magnetic field in the Snake suggesting a strong connection between the physics of the anchoring molecular regions and the Snake.  We suggest that the physical processes considered here may be relevant to many of the radio filaments near the Galactic Center. We also suggest further observations of the Snake and other filaments that would be useful for obtaining further insights into the physics of these objects. ",
        "introduction": "The fundamental nature of the numerous radio filaments observed near the Galactic Center (see \\citet{larosa00a} for an overview) is unclear. In particular: How are they formed? What are the sources of the relativistic electrons? Why do they have such large magnetic field strengths? What is the reason for the unusually flat spectral indices in some of them? The discovery of the Snake \\citep{gray91a}, a 60 pc by 0.4 pc  filament, approximately 150 pc to the West of Sgr A, resulted in the additional feature that the characteristics of the radio emission are strongly related to a morphological  ``kink''. Further observations \\citep{gray95a}, revealed a major and a minor kink, with the radio intensity and spectral index systematically varying away from the major kink. Various models have been proposed to explain the spectacular elongated structure of the Snake (see \\citet{gray95a}). However, none of them have been able to successfully explain its major features. In our view, the Snake is in many respects one of the least complicated of the filamentary features close to the Galactic Center and a good physical model may hold the key to understanding the numerous arcs and filaments in the Galactic Center region.  We propose that the Snake is a magnetic flux tube with both ends anchored in dense rotating material (molecular clouds and/or associated HII regions). Differential rotation of the tube at either or both ends produces a monotonically increasing toroidal magnetic field and when this reaches a critical value ``coils'' or ``loops'' are formed through localized kink instabilities. Release of the magnetic energy stored in a coil, through magnetic reconnection (and possibly shocks), is the source of the relativistic electron energy. Energetic electrons diffuse away from each coil at an energy-dependent rate causing a flattening spectral index. ",
        "conclusions": "High resolution images of the Snake near the major kink (\\citet{gray95a}, Figures 10 and 11) show that the Snake appears to be split in two, similar to flux tubes in the late stages of the \\citet{bazdenkov98a} simulations. This provides additional support for our proposal. As well as being directly applicable to the Snake, our model may open up a number of appealing possibilities for the dynamics of other magnetic filaments in the interstellar medium near the Galactic Center. It is feasible, in an environment with such a large density of molecular clouds, that rotation of  magnetic flux tubes threading their cores would lead to  kink instabilities and reconnection. The flattening spectral index of the Snake depends upon the differential diffusion of energetic particles and this is such a well known phenomenon in the regular ISM that its role in this model is unremarkable.  The index of the momentum dependence of the diffusion parameter that we have inferred for the Snake is in the range of values used in models of cosmic ray propagation. More detailed models raise the prospect of a better determination of this index as well as the other parameters in this model. In particular, a more precise value of the  index of the momentum dependence of the particle creation rate, which our current estimates place near the value associated with strong shocks, would be of interest. The type of data that is essential for more detailed modeling include (1) High resolution images showing more clearly the structure of the flux tube over its entire length. (2) Flux densities at different frequencies, enabling one to better constrain the energy dependence of the diffusion parameter and the magnetic field.  It remains to be shown that flux tubes in the Galactic Center can be wound up to such an extent that they produce a significant toroidal field. If an anchoring cloud acquires an angular velocity which persists for long enough that it rotates through 300 radians ($\\sim 50$ revolutions), that would be sufficient. Another possibility is that the required winding could be produced in a contracting cloud, although this idea requires further investigation. In our independent estimates of the magnetic field in the Snake and in the related issue of subcritical or supercritical contraction we have concentrated exclusively on the molecular complex at the northern end of the Snake. Currently, however, there is little known about the southern end and the masses and rotation rates of molecular clouds in that region will certainly be of interest."
    },
    "0011/astro-ph0011218_arXiv.txt": {
        "abstract": "VLT spectra of 14 luminous compact galaxies (LCGs) reveal strong metallic absorption line sytems as well as narrow and intense emission lines. Their gas extinction is found to be large (A$_V$ $\\sim$ 1.5 mag) leading to an upward revision of their star formation rate (SFR) to an average value of $\\sim$ 40 $M_{\\odot} yr^{-1}$. Large extinction values are also supported by the large rate of detection in one field observed by ISO. Gas metal abundances in LCGs have about half the solar value. LCG absorption spectra can be synthesized with a mix of a few Gyr old and relatively metal-rich (generally solar to over-solar values) stellar population and a younger stellar population ($<$ 5$\\times 10^{8}$ years) having a metal abundance similar to that of the gas. \\\\  We argue that LCGs are the progenitors of present-day spiral bulges. LCGs have masses and light concentrations similar to those of present-day bulges. They could have been formed entirely during a period of a few Gyr prior to the epoch of their observations if the star formation has been sustained at the observed rate. As in present-day galactic bulges, LCG stars show a wide range of abundances. Thus, observing LCGs allows us to directly witness an important stage in the formation of a massive galaxy, the building of the bulge prior to that of the disk. The gas needed to feed the observed star formation is likely to be falling in from the outskirts of the galaxy, being tidally pulled out from interacting companion galaxies. An infall scenario naturally explains the gas metal abundance which is generally lower than that of the older stellar component. At least for the strongest star-forming LCGs, there is clear imaging evidence for the presence of companions. Some LCGs also show evidence for the beginning of a disk formation. \\\\ If the above scenario holds for most LCGs, we estimate that at least 20\\% of present-day spiral galaxies have formed the bulk of their stars at relatively recent epochs, during the last 8-9 Gyr, at redshifts less than $\\sim$ 1. Since they are heavily extincted, we predict their IR luminosities to be relatively large, around $L_{IR} $=$ 10^{11}$ $L_{\\odot}$, i.e. near or slightly below the luminosities of the galaxies detected by ISO in the same redshift range. Taking into account the integrated IR luminosity of the LCG galaxy population can lead to a significant upward revision of the cosmic SFR density in the redshift range from 0.5 to 1. ",
        "introduction": "The evolution of UV and IR luminosity densities can be understood as a decrease by a factor of $\\sim$ 6 of the global star formation density from the epoch corresponding to redshift z=1 to the present epoch (Lilly et al. 1996; Flores et al. 1999). These luminosity densities evolve similarly and contribute equally to the star formation in that redshift range, if one assumes that the shapes of the UV and IR luminosity functions to be similar to the local ones (Hammer 1999). Integration of the derived global star formation implies that between half and two thirds of the present-day stars have been formed since z=1. This is somewhat in contradiction with the primordial collapse scenario in which the most massive galaxies were formed at much earlier epochs, at redshifts larger than $\\sim$ 3 (e.g. Renzini and Cimatti 1999).  The latter scenario is supported by the apparent non-evolution in the number density of large disks (Lilly et al. 1998), and by its success in reproducing galactic disk evolution with chemical evolution models normalized to local galaxies (Boissier and Prantzos 2000). \\\\ The study of the luminosity density evolution relies on observations of luminous galaxies with $M_{B} <$ -20 in the optical, and with $L_{bol} >$ 2 $\\times$ $10^{11}$ $L_{\\odot}$ in the IR.  The two main galaxy populations responsible for the observed decline in the star formation density are, in decreasing importance, the luminous IR galaxies (Flores et al. 1999) and the luminous compact galaxies (LCG) (Guzman et al. 1997; Lilly et al. 1998).  The former are generally large and massive galaxies found mostly in interacting systems (Hammer, 1999), and which show star formation rates (SFR) exceeding 50 $M_{\\odot}yr^{-1}$. They represent $\\sim$ 4\\% of the field galaxy population.  The latter are much more numerous and contribute to more than 40\\% of the UV luminosity density at z$\\sim$ 1 (Guzman et al. 1997; Hammer \\& Flores 1998). They constitute the most rapidly evolving galaxy population seen in the UV.  The properties of these LCGs are still largely unknown. They are sometimes so compact that their nucleus is barely resolved in the WFPC2 HST images, which prevents their classification on the basis of their luminosity profile.\\\\ \\\\  Koo \\& Kron (1981) first drew attention to these compact galaxies during their search for faint QSOs on the basis of colors and compactness.  They identified a sparse population (the surface density is $\\sim$ 30 sources per square degree) of very blue compact galaxies with z $\\le$ 0.7 and exhibiting narrow emission lines.  Using the Keck HIRES spectrograph, Koo et al. (1995) demonstrated that the emission lines of LCGs have a Gaussian profile, and suggested that they may be progenitors of local spheroidal galaxies through fading by $\\sim$ 4 to 5 magnitudes. More recently, Guzman et al. (1997) and Phillips et al. (1997) have identified a large population of emission-line compact galaxies in the flanking fields around the Hubble Deep Field (HDF, Williams et al.  1996) at z $\\le$ 1, comprising $\\sim$ 20\\% of the field galaxy population.  They found the [OII]$\\lambda$3727 emission-line velocity widths of the compact galaxies to range between 35 and 150 km s$^{-1}$ (Phillips et al. 1997), comparable to those of local dwarf galaxies, although the compact galaxies are 10 to 100 hundred times more luminous (up to a few $L^{\\star}$). The similarities between the HDF compact galaxies and those from Koo and Kron (1981) have led Guzman (1999) to suggest that blue compact galaxies at intermediate redshifts may be the progenitors of local spheroidal or irregular galaxies through fading by $\\le$ 4 magnitudes. If true, according to Guzman (1999), these should have extremely low M/L ratio, and this scenario would imply a severe downwards revision of the amount of star formation at z$\\sim$ 1, since the baryonic content in irregulars (or spheroidals) is very small (Fukugita et al. 1998).\\\\ \\\\  We develop in this paper a different scenario for the nature of the brightest LCGs, which are major contributors to the observed decline of star formation density since z=1.  In our picture, LCGs are not the progenitors of present-day spheroidal and irregular galaxies, but rather of today's spiral bulges which built up gradually over time by a series of mergers of smaller units.  There have been several previous observations that have hinted at that scenario, but their meaning was not realized at the time. Thus, Phillips and al. (1997) found that a large fraction of galaxies selected by their compactness appeared to be small but bright spirals.  Even adding a color criterion is not sufficient to distinguish between old and young stellar populations: Schade et al. (1999) found that fully a third of the distant ellipticals they studied show star formation and colors as blue as those of blue compact galaxies.  The merging scenario comes naturally as the merging rate was about ten times higher at z=1 than it is today (Le F\\`evre et al. 2000). Moreover, Guzman (1999) and Hammer (1999) have noticed that a large fraction of compact galaxies selected from the HDF and Canada-France redshift survey (CFRS) fields are in interacting systems, with the presence of complex and distorted features surrounding their cores.  Finally, the LCGs in the CFRS fields have near-IR and visible colors very similar to those of other luminous non-dwarf galaxies, with values apparently not consistent with those produced by a strong and short burst of star formation in a low-mass object (Hammer et al. 1997).  We present next new observational data which support and sharpen the above scenario.\\\\ In the following, $h_{50}$=$H_{0}$/50 and if not specified we adopt $h_{50}$=1 and $q_{0}$=0.5. ",
        "conclusions": "Spectroscopic observations of a small sample of 14 Luminous Compact Galaxies (LCG) with the VLT have provided new and important information on a crucial stage in the formation of galaxies, that of bulges in massive (non-dwarf) spiral galaxies. The LCG spectra are characterized by the following properties:  1) they show strong metallic absorption lines, including those from $\\alpha$-elements and iron.  2) they show large Balmer decrement ratios (as measured by $H\\beta/H\\gamma$), implying large extinctions (median $A_V$ of $\\sim$ 1.5 mag).  3) the SFRs derived from the extinction-corrected Balmer lines are more than ten times higher than the corresponding ones derived from the UV luminosities.  LCGs with redshifts between 0.4 and 0.8 are undergoing intense bursts of star formation, with SFRs averaging to $\\sim$ 40 $M_{\\odot} yr^{-1}$.  Because of their large extinctions, they are expected to be luminous in the IR, at a luminosity level just below the IR luminous galaxies detected by ISO in the same redshift range (Flores et al. 1999).  Indeed 3 of the 14 LCGs discussed in this paper are detected by ISO, implying higher extinction for them.  We believe that LCGs are progenitors of present-day bulges of massive spiral galaxies because:  1) they have stellar masses similar to those of today's bulges in $\\ge$$L^{\\star}$ spirals, concentrated in comparable volumes.  2) their stars show a wide range of metal abundance, from sub-solar to over-solar values, as in the bulge of the Milky Way.  3) the majority of LCGs show low-surface-brightness components around their high-surface-brightness cores (the bulges). Most are in interacting systems, and the tidally pulled gas from the companions is likely to sustain star formation in the LCGs at a high rate for several hundreds Myr, resulting in the completion of the bulge formation and the beginning of the disk formation.  The characteristic time for forming the bulk of the stars in LCGs is rather small, only $\\sim$ 1.6 Gyr on average. It is comparable to the age of the oldest stars in LCGs derived from the analysis of their absorption lines. This means that a significant fraction (several tens of percent) of present day's massive galaxies could have been formed at relatively recent epochs, at z $\\le$ 1.  This is consistent with the cosmic star formation history which predicts that an important fraction of the stellar mass was formed in the z = 0 -- 1 redshift range (Hammer 1999).  The importance role played by galaxy interactions in the formation of bulges of spiral galaxies is also consistent with the hierarchical theory of galaxy formation. On the other hand, the scenario outlined here may appear to be in contradiction with the results of Lilly et al. (1998) who found that the number density of large spirals at z $\\sim$ 1 is comparable to that today. However the Lilly et al. results are based on only $\\sim$ 40 morphologically selected large spirals in the redshift range from 0.2 to 1, and their large one $\\sigma$ error bar can easily accomodate a density evolution of spiral disks of $\\sim$ 20\\% as found here.  Moreover, some of the large spirals in Lilly et al. (1998)' sample are in interacting systems and are strong infrared emittors, such as CFRS 03.1540 discussed here, or other examples discussed by Flores et al. (1999). Presumably, these large interacting disks will merge leading to a density evolution.  If the 14 LCGs described here are representative of the whole LCG population at z $> $0.4, then their contribution to the global star formation density should be revised upwards by factors as large as 7 -- 10. That contribution could then be as high or higher than that of the luminous IR galaxies detected by ISO (Flores et al. 1999). The latter are more massive systems and formed stars more rapidly than LCGs, but they constitute a much sparser population of galaxies, comprising only a few percent of the total population. Most LCGs at z $\\sim$ 0.5 would have bolometric IR luminosities not far below the ISO sensitivity limit, and we predict that they will be easily detected in large numbers by SIRTF.  Finally we reiterate that our interpretation of the nature of LCGs is quite different from that proposed by Koo et al. (1995), Guzman et al. (1997) and Guzman (1999) for the majority of compact galaxies found in the Hubble Deep Field and flanking fields. This is probably because most of the Guzman et al galaxies have intrinsically lower luminosities than LCGs.  We believe LCGs to be 10 to 100 times more massive than dwarf galaxies, to be forming stars at large rates, so they cannot be the progenitors of local spheroidal or irregular systems through fading. Corresponding to the brightest 25\\% of the compact galaxies selected by Guzman et al (1997), the LCGs selected in CFRS fields are mostly evolved starbursts probably similar to those classified as such in the HDF. Due to their large derived SFRs, LCGs are the main contributors to the large SFR density observed at z=0.5 to 1.  In future papers (Gruel et al. in preparation) we will investigate the velocity fields and dynamics of LCGs, using data of superior quality and based on a larger sample. We will also be studying the possible revision of the star formation history and of the IR luminosity density caused by the LCG population."
    },
    "0011/astro-ph0011504_arXiv.txt": {
        "abstract": "We construct numerical models of mildly triaxial elliptical galaxies with central density cusps. Using a technique we call ``adiabatic squeezing,'' we begin with a spherical $\\gamma=1$ Hernquist model and apply a drag to the velocities of the particles along each principle axis. The final models are stable in isolation, preserving their density structure and figure shape over many dynamical timescales. The density profile and axial ratios compare well to the observed properties of elliptical galaxies. The orbital structure of these models show a mixture of tubes, boxes, and boxlets, as expected for triaxial systems, with very few chaotic orbits. These $N$-body realizations of cuspy triaxial galaxies provide a basis for the study of dynamical evolution of elliptical galaxies. ",
        "introduction": "There is good reason to believe that many elliptical galaxies are at least moderately triaxial systems. Observationally, the distribution of isophotal shapes of ellipticals argues that there are few perfectly spherical or axisymmetric members of the present-day elliptical population (Tremblay \\& Merritt 1995; Ryden 1996). Using morphological and kinematic data, Franx, Illingworth, \\& de Zeeuw (1991) have shown that ellipticals show mild triaxiality, while more recently Bak \\& Statler (2000) have arrived at a similar conclusion by modeling the kinematics of ellipticals.  There is strong theoretical support for triaxiality as well. Under the gravitational collapse model for elliptical galaxy formation, models of protogalaxies undergo a non-spherical collapse which results in a virialized system that is decidedly triaxial (Dubinski \\& Carlberg, 1991). Similarly, in the hierarchical merging picture, the remnants of simulated galaxy mergers also show strong triaxiality (\\eg Hernquist 1992; Barnes 1988, 1992; Hernquist 1993), even when multiple mergers are considered (Weil \\& Hernquist 1996). Either model of elliptical galaxy formation predicts that a dynamically young elliptical should possess a fair degree of triaxiality, although subsequent evolution may drive the system towards axisymmetry.  There is also good reason to believe that many ellipticals have cuspy density profiles. While elliptical galaxies were once characterized by flat, isothermal cores, studies using the {\\it Hubble Space Telescope} have shown that in fact the density profile of ellipticals rarely becomes perfectly flat in the inner region. Instead, the density profile continues to rise as $\\rho \\propto r^{-\\gamma}$, where $0.25 < \\gamma < 2$ (\\eg Lauer \\etal 1995; Byun \\etal 1996; Gebhardt \\etal 1996; Faber \\etal 1997).  Furthermore, correlations exist between the cusp slope $\\gamma$ and the global properties of the galaxy, including shape in the form of boxy or disky isophotes (Faber \\etal 1997). Dynamical arguments suggest also that the presence of a strong central cusp ($\\gamma > 1$) acts to drive chaos in the orbit families which populate the galaxy, driving the system away from strong triaxiality (\\eg Gerhard \\& Binney 1985; Norman, May, \\& van Albada 1985; Merritt \\& Valluri 1996; Holley-Bockelmann \\etal 2000).  While real elliptical galaxies most probably are cuspy and moderately triaxial, most self-consistent modeling of ellipticals to date has focussed either on cuspy spherical or axisymmetric models, or triaxial models with constant density cores (\\ie $\\gamma=0$). The perfect ellipsoid of de Zeeuw (1985) allows for differing degrees of triaxiality, but has a flat $\\gamma = 0$ core.  Self-consistent $N$-body models of elliptical galaxies used in studies of black hole growth have either examined triaxiality in a $\\gamma \\sim 0$ profile (\\eg Norman \\etal 1985; Merritt \\& Quinlan 1998), or focused on cuspy models which are spherical (Sigurdsson \\etal 1995; Quinlan \\& Hernquist, 1997) or axisymmetric (van der Marel \\etal 1997). The problem lies in the lack of an analytic distribution function for systems which are both cuspy and triaxial; without such a distribution function, generating initial conditions for simulations is difficult.  Here we employ an alternative approach to generating triaxial galaxy models with well-defined shapes and cusps. We apply an adiabatic drag on the particles in a spherical $\\gamma=1$ Hernquist model to mold the system into triaxial shape. This process preserves the cusp slope and results in models which are moderately triaxial and have axis ratios that are reasonably constant with radius. More importantly, the models are stable, exhibiting little evolution in their structural properties over many crossing times. We characterize the orbit families which populate such models, and find a rich mixture of tubes, boxes, and boxlets, with very few chaotic orbits. This technique has also been applied successfully in a $\\gamma = 0$ model, and can, in principle, generate models with a wide range in cusp slopes. These well-defined models thus present a useful tool for studying the evolution of elliptical galaxies or triaxial halos. ",
        "conclusions": "Using a technique we refer to as ``adiabatic squeezing'' we have constructed models of triaxial galaxies with central density cusps. These models preserve a Hernquist (1990) $\\gamma=1$ profile over many orders of magnitude, from $-2 < \\log(r/a) < 2$, or 10 pc $ < r < $ 100 kpc when scaled to a luminous elliptical with half mass radius 1.5 kpc. The models possess fairly constant axis ratios as a function of ellipsoidal radius, showing the range of triaxiality ($0.3 < T < 0.75$) and modest flattening inferred for massive ellipticals. The models are stable in isolation, showing little evolution in either the density profile or figure shape over many dynamical times.  A combination of fourier spectral classification and axis-crossing pattern recognition have been used to classify the orbit content of a characteristic model. At the most tightly bound energies, the models possess a rich variety of orbits, including tubes, boxes, and resonant boxlets, reflecting the triaxial nature and central cusp of the model. Moving to progressively less bound energy slices, the ratio of loops:boxes rises such that the outer regions contain predominantly loop orbits. While the presence of noise in the potential expansion limits our ability to detect modest chaos in the orbit populations, we find no strongly chaotic orbits in the model of the type that would result in rapid evolution of the model.  These models thus represent a useful tool for studying the self-consistent dynamical evolution of triaxial galaxies and halos. We are currently using these techniques to explore the evolution of elliptical galaxies harboring central massive black holes (Holley-Bockelmann \\etal 2001). Other potential uses include studying the inflow of gas in triaxial systems, the response of triaxial dark halos to baryonic accretion and disk formation, and the infall of satellite companions into elliptical galaxies."
    },
    "0011/astro-ph0011397_arXiv.txt": {
        "abstract": "We examine the three--dimensional clustering of \\CIV\\ absorption--line systems, using an extensive catalog of QSO heavy--element absorbers drawn from the literature. We measure clustering by a volume--weighted integral of the correlation function called the reduced second--moment measure, and include information from both {\\em along} and {\\em across} QSO lines of sight, thus enabling a full determination of the three--dimensional clustering of absorbers, as well as a comparison of line-- and cross--line--of--sight clustering properties. Here we present the three--dimensional reduced second--moment estimator for a three--dimensional point process probed by one--dimensional lines of sight, and apply our algorithm to a sample of 345 \\CIV\\ absorbers with median redshift $\\langle z\\rangle = 2.2$, drawn from the spectra of 276 QSOs. We confirm the existence of significant clustering on comoving scales up to 100 \\hMpc\\ ($q_0 = 0.5$), and find that the additional {\\em cross}--line--of--sight information strengthens the evidence for clustering on scales from 100 \\hMpc\\ to 150 \\hMpc. There is no evidence of absorber clustering along or across lines of sight for scales from 150 \\hMpc\\ to 300 \\hMpc. We show that with a 300--times larger catalog, such as that to be compiled by the Sloan Digital Sky Survey (100,000 QSOs), use of the full three--dimensional estimator and cross--line--of--sight information will {\\em substantially} increase clustering sensitivity. We find that standard errors are reduced by a factor 2 to 20 on scales of 30 to 200 \\hMpc, {\\em in addition} to the factor of $\\sqrt{300}$ reduction from the larger sample size, effectively increasing the sample size by an extra factor of 4 to 400 at large distances. ",
        "introduction": "In a previous series of investigations \\citep{VQYY96,Quash96,Quash98}, the clustering properties of \\CIV\\ and \\MgII\\ absorbers have been investigated, using an extensive catalog of heavy--element absorption--line systems drawn from the literature.\\footnote{Contact D. E. Vanden Berk (danvb@fnal.gov) for the latest version of the catalog; see \\citet{Y91} for an earlier version.} These authors used a one--dimensional correlation analysis --- one confined to pairs of absorbers along the {\\em same} QSO line of sight --- and found  evidence for strong and evolving clustering on small scales (1--16 \\hMpc), as well as for superclustering on scales as large as 50--100 \\hMpc. Together, these investigations suggest that these strong absorbers, with median rest equivalent width $\\langle W\\rangle=0.4$ \\AA\\ for \\CIV\\ \\citep{Quash98}, are biased tracers of the higher density regions of space, and that agglomerations of absorbers along a line of sight are indicators of clusters and superclusters.  Recently, \\citet{Quash99} used a new measure of clustering, called the reduced second moment measure or $K(r)$ \\citep{Ripley88,Badd98}, which directly measures the mean overdensity of absorbers on scales $\\la r$. While closely related to other second--order measures of clustering, such as the correlation function or the power spectrum, the reduced second moment measure nevertheless has a number of advantageous statistical properties, and has recently been studied by astrophysicists \\citep{Mart98,Quash99,Stein00}. \\citet{Quash99} found significant evidence for clustering of \\CIV\\ absorbers on scales of 20--100 \\hMpc, with marginal evidence on 100--200 \\hMpc\\ scales, again suggesting that the absorbers show superclustering much like what is seen locally in the distribution of galaxies.  However, their analysis also was confined to pairs of absorbers along the same line of sight, and hence did not include {\\em cross}--line--of--sight information valuable for a full determination of the three--dimensional clustering properties of absorbers, which we are ultimately interested in finding. Indeed, \\citet{RY99} have claimed that there is evidence of some significant contamination of true intervening systems along the line of sight by absorbers that are actually physically associated with the QSO, and that such contamination may extend to relative velocities as great as 75000 \\kms. This means that such a contamination could be present as well in the superclustering signal found by \\citet{Quash99}. In addition, on smaller scales of a few megaparsecs, \\citet{Crotts97} have claimed that the clustering of \\CIV\\ systems {\\em across} adjacent lines of sight is significantly weaker than {\\em along} a line of sight and have questioned whether the correlations are due to velocity dispersion in associated systems rather than intrinsic spatial clustering. The number of adjacent lines analyzed by \\citet{Crotts97} was small, however.  It is thus essential to study correlations both {\\em along} and {\\em across} QSO lines of sight, to enable a complete determination of the three--dimensional clustering of absorbers, as well as a comparison of line-- and cross--line--of--sight clustering properties (and ultimately statistically discriminate between intrinsic and intervening absorbers). What is required is a method that measures spatial clustering of absorbers --- which are nonetheless confined to QSO lines of sight --- and can contrast clustering along and across lines of sight.  Here we present the three--dimensional estimator for the reduced second moment measure of a three--dimensional point process probed by one--dimensional lines of sight. We apply our algorithm to a sample of 345 \\CIV\\ absorbers with median redshift $\\langle z\\rangle = 2.2$, drawn from the spectra of 276 QSOs in the aforementioned catalog of Vanden Berk \\etal, with the goal of determining the clustering of absorbers on very large scales and seeing if the signal found by \\citet{Quash99} remains when including cross--line--of--sight information, or, more compellingly, using only cross--line--of--sight information.  The outline of the paper is as follows: In \\S 2 we define the three--dimensional reduced second moment measure, compare it to its one--dimensional analog, and outline our method of estimation. In \\S 3 we apply our methodology to the above sample and present our results. In \\S 4 we interpret the results and discuss how the full three--dimensional estimator and cross--line--of--sight information will improve our ability to measure clustering with the absorber sample from the Sloan Digital Sky Survey (100,000 QSOs). Finally, we conclude in \\S 5 and present the details of the three--dimensional reduced second moment estimator in the Appendix. ",
        "conclusions": "We present two estimators, $\\hat{K}_\\parallel(r)$ and $\\hat{K}(r)$, for the three--dimensional reduced second moment for one--dimensional data (absorber redshifts) along QSO lines of sight. The first estimator uses absorber pairs along the {\\em same} lines of sight, whereas the latter includes data from across {\\em different} lines of sight. We apply our algorithm to a sample of 345 \\CIV\\ absorbers with median redshift $\\langle z\\rangle = 2.2$, from the spectra of 276 QSOs, drawn from the catalog of Vanden Berk \\etal.  We confirm the existence of significant clustering of \\CIV\\ absorbers on comoving scales up to 100 \\hMpc\\ ($q_0 = 0.5$), and find that the additional {\\em cross}--line--of--sight information strengthens the evidence for clustering on scales from 100 \\hMpc\\ to 150 \\hMpc. This argues against claims that all the apparent clustering on these scales is due to significant contamination along the line of sight by absorbers that are actually physically associated with the QSO. However, the limited size of that catalog, as well as its heterogeneity, precludes a final, strong statement of the statistical significance and amplitude of the clustering on scales $\\ga 100$ \\hMpc. Also, the sample is too sparse to significantly compare clustering along and across lines of sight on scales of less than 100 \\hMpc. There is no evidence of absorber clustering along or across lines of sight for scales from 150 \\hMpc\\ to 300 \\hMpc.  We show that with a 300--times larger catalog, such as that to be compiled by the Sloan Digital Sky Survey (100,000 QSOs), use of the full three--dimensional estimator and cross--line--of--sight information will {\\em substantially} increase clustering sensitivity. We find that standard errors are reduced by a factor 2 to 20 on scales of 30 to 200 \\hMpc, {\\em in addition} to the factor of $\\sqrt{300}$ reduction from the larger sample size, effectively increasing the sample size by an extra factor of 4 to 400 at large distances. Thus, use of the full three--dimensional reduced second moment estimator will significantly advance our ability to describe and analyze large--scale clustering of absorbers, and hence visible matter, from the SDSS Catalog."
    },
    "0011/astro-ph0011168_arXiv.txt": {
        "abstract": "We present new redshifts and positions for 635 galaxies in nine rich clusters containing Wide-Angle Tailed (WAT) radio galaxies. Combined with existing data, we now have a sample of 18 WAT-containing clusters with more than 10 redshifts. This sample contains a substantial portion of the WAT clusters in the VLA 20 cm survey of Abell clusters, including 75\\% of WAT clusters in the complete survey (z$\\leq$0.09), and 20\\% of WAT clusters with z$>$0.09. It is a representative sample which should not contain biases other than selection by radio morphology. We graphically present the new data using histograms and sky maps. A semi-automated procedure is used to search for emission lines in the spectra in order to add and verify galaxy redshifts. We find that the average {\\em apparent} fraction of emission line galaxies is about 9\\% in both the clusters and the field. We investigate the magnitude completeness of our redshift surveys with CCD data for a test case, Abell 690.  This case indicates that our galaxy target lists are deeper than the detection limit of a typical MX exposure, and they are 82\\% complete down to R=19.0. The importance of the uniformity of the placement of fibers on targets is posited, and we evaluate this in our datasets.  We find some cases of non-uniformities which may influence dynamical analyses. A second paper will use this database to look for correlations between the WAT radio morphology and the cluster's dynamical state. ",
        "introduction": "Wide-angle tailed radio galaxies (WATs) are a powerful type of FR-I (``edge-darkened\", \\cite{FR74}) radio galaxy associated with first-ranked ellipticals in clusters. Generally, the tails of a WAT bend in a common direction, giving the source an overall {\\sc V}, or {\\sc C} shape (see   % O'Donoghue \\etal 1990 for a radio survey). The mechanism which bends the WATs is poorly understood. Ram pressure caused by motion through the intracluster medium (ICM) was suspected from the time of discovery (\\cite{OR76}). The sticking point in this hypothesis is that \\app 1000 \\kms\\ relative velocities may be required for bending (\\cite{Eil84}), and the host galaxies, typically gE, D and cD, should have much smaller peculiar velocities with respect to their cluster (\\cite{Mal92}, \\cite{Bu86}).  A subcluster merger hypothesis has emerged linking the bent morphology of WATs to the dynamical state of the host cluster (Pinkney \\etal 1993, hereafter \\cite{Pi93}, Gomez et al 1997, \\cite{Rot96}, \\cite{Lok95}). In it, the WAT acts as a stationary weathervane, its tails indicating the local flow of ICM in the stormy environment of a cluster-subcluster merger (\\cite{Burns98}). If this hypothesis is true, then WAT's should be beacons of cluster mergers.  This hypothesis is difficult to confirm for a {\\em single} cluster: only {\\em radial} peculiar velocities can be ruled out for the WAT galaxy, and the optical signatures of merger can be subtle during post-merger epochs when the ICM motion is still sufficient to shape the WAT (\\cite{Lok95}; \\cite{Pi96}). A large database can provide statistical leverage to test the subcluster-merger hypothesis. In particular, a distribution of WAT radial peculiar velocities can constrain the true WAT velocities, thereby testing the moving-WAT hypothesis. Moreover, the occurrence rate of significant substructure can be compared to a control sample of radio-quiet clusters. In this paper, we present new velocities and positions of over 600 galaxies in 9 WAT clusters. This will be combined with published data for 9 additional WAT clusters. These data will be used to probe the connection between WAT bending and cluster evolution in paper II. In this paper, section 2 will contain the sample selection criteria. Section 3 describes the observations and reductions. In section 4, we compare our velocities to published ones and discuss the emission line content. Finally, in section 5 we examine the completeness and spatial uniformity of our redshift surveys. ",
        "conclusions": "We have presented new redshifts and positions for galaxies in 9 WAT clusters. Nineteen of our galaxies had published redshifts which are in reasonable agreement with our redshifts (i.e., $\\sigma \\simlt 150$). A search for emission line galaxies helped contribute cluster redshifts and improve the accuracy of our catalog. Our candidate galaxies for observation were chosen by visual inspection of the PSS. CCD data for Abell 690 indicate that its candidate list ({\\it MX file}) is 82\\% complete down to R$\\sim$19.0. Also, 95\\% of A690 targets are observed, while about 60\\% of them have reliable redshifts down to the MX detection limit (R$\\sim$18.5). The fraction of candidates that are {\\em observed} varies from cluster to cluster, and, for 2 clusters, varies strongly between subregions in the cluster. Nevertheless, the dataset is well-suited for measuring WAT peculiar velocities, and can also be used for substructure analyses with attention to the cases with uneven sampling. Paper II will use these data to look for links between the WAT radio morphology and the cluster's dynamical state."
    },
    "0011/astro-ph0011442_arXiv.txt": {
        "abstract": "{RRd stars from the Magellanic Clouds form a well-defined band in the Petersen diagram. We explain this observed band with our evolutionary and pulsation calculations with assumed metallicity [Fe/H]$=(-2,-1.3)$. Vast majority of RRd stars from LMC is confined to a narrower range of $(-1.7,-1.3)$. The width of the band, at specified fundamental mode period, may be explained by mass spread at given metallicity. The shape of the band reflects the path of RRd stars within the RR~Lyrae instability strip. We regard the success in explaining the Petersen diagram a support for our evolutionary models, which yield, mean absolute magnitude in the mid of the instability strip, $\\left<M_V\\right>$, in the range $0.4$ to $0.65$ mag implying distance modulus to LMC of $18.4$ mag.}{stars: variable, stars: oscillations, stars: RR~Lyrae, stars: double-mode pulsations, galaxies: Magellanic Clouds, stars: abundances} ",
        "introduction": "RR~Lyrae stars are objects of great importance for whole astrophysics. Double-mode RR~Lyrae stars (RRd), constitute relatively rare subtype. Nonetheless RRd stars attract considerable attention because they provide us an additional precise observable, which is the second period (see review of Kov{\\'a}cs (2000a)).  Usefulness of double-mode pulsators was first realized by Petersen~(1973). Following his idea, the period data for double-mode pulsators are commonly represented in diagrams, called now {\\em Petersen diagrams}, in which the period ratio, ${\\cal R}=P_1/P_0$, is plotted against the fundamental mode period, $P_0$. Published examples of Petersen diagrams for RRd stars may be found in the following papers: Nemec~\\etal (1985a), Clement~\\etal (1986), Walker~\\etal (1994), Alcock~\\etal (1997,2000b), Beaulieu~\\etal (1997). The Oosterhoff dichotomy is manifested in Petersen diagrams for RRd stars. Namely, RRd variables from the Oosterhoff~I globular clusters occur in systematically shorter period range than those from the Oosterhoff~II (Smith 1995).  Petersen diagram in its first application was used to derive mass and radius estimates for Cepheids (Petersen 1978). The result was a huge discrepancy, known as the double-mode Cepheid mass problem, between the masses derived in this way and the evolutionary masses. The problem disappeared once the OPAL opacities became available (Moskalik~\\etal 1992).  First comparison of the observed and theoretical Petersen diagrams for RR~Lyrae stars was made by Cox~\\etal (1980). No large discrepancy between observations and models was found. The new opacities were first used to calculate period ratios for RRd models by Cox (1991), later by Kov\\'{a}cs~\\etal (1991,1992) and Bono~\\etal (1996). Kov\\'{a}cs~\\& Walker (1999) used data on RRd stars from globular clusters to derive luminosities of RR~Lyrae stars yielding support for a brighter luminosity scale.  A large number of data on RRd stars became available as a byproduct of micro-lensing surveys. Most of currently known RRd stars were observed by MACHO collaboration (Alcock~\\etal 1997,2000b). Figure~1. shows Petersen diagrams for all known RRd stars. The objects belonging to different stellar systems, galaxies and globular clusters, are shown in separate panels with big dots. We see that RRd stars form quite a narrow curved band in the Petersen diagram. The objects from LMC are spread all over the period range, which is $0.46-0.6$ d. Short period end in LMC is significantly more populated. We don't see such a concentration for the SMC objects which are more-or-less uniformly distributed, however in view of much scarcer data this conclusion must be regarded preliminary. The data for the Galaxy and dwarf galaxies Draco and Sculptor are too sparse to conclude anything about properties of RRd stars in these systems. The objects from globular clusters are localized in small parts of the band.  The goal of our paper is to explain the properties of the RRd band as defined by the objects from LMC and SMC. From this we expect to learn something about RR~Lyrae star properties, particularly their luminosities, as well as about early history of star formation in the Magellanic Clouds. Furthermore, the abundant data on RRd stars may yield us a useful constraint regarding the origin of double-mode pulsation. Although the double-mode pulsation has been successfully modeled (e.g. by Feuchtinger (1998)), our understanding of its causes is not yet satisfactory.  Application of Petersen diagrams as a probe of stellar properties is explained in Section~2. In Sections~3. and~4. we provide some details regarding our model and pulsation calculations, respectively. Interpretation of the RRd band based on model calculation are presented in Section~5. In Section~6. we discuss absolute magnitudes of RRd stars. Uncertainties of the theoretical Petersen diagram are analyzed in Section~7.  \\begin{figure}[hbt] \\centering \\epsfig{file=fig1.ps,height=1.27\\linewidth,bbllx=21,bblly=40,bburx=572, bbury=740} \\caption[]{{\\em Petersen diagrams} for RRd stars in various stellar systems. Data taken from (number of objects in parentheses): {\\bf LMC}~(181) -- Alcock~\\etal (1997,~2000b), {\\bf SMC}~(26) -- OGLE private communication (unpublished), {\\bf Draco}~(10) -- Nemec (1985a), {\\bf IC4499}~(16) -- Clement~\\etal (1986), Walker~\\&~Nemec (1996), {\\bf M3}~(5) -- Corwin~\\etal (1999), {\\bf M15}~(12) -- Nemec (1985b), Jurcsik~\\&~Barlai (1990), Purdue~\\etal (1995), {\\bf M68}~(12) -- Walker (1994), {\\bf Sculptor}~(1) -- Ka{\\l}u{\\.z}ny~\\etal (1995), {\\bf Galactic field}~(5) -- Garcia-Melendo~\\&~Clement (1997), Clementini~\\etal (2000).} \\label{f:obs} \\end{figure} ",
        "conclusions": "\\label{s:sum}  We have seen that the Petersen diagram for RRd stars in the Magellanic Clouds may be explained by standard model calculations. The calculations involved evolutionary models of horizontal branch stars and linear non-adiabatic calculations of radial pulsation periods. The agreement has been achieved assuming metal abundances consistent with other determinations.  This successful explanation of the Petersen diagram may be regarded as a test of our models. These models yield mean absolute magnitudes, $\\left<M_V\\right>\\approx0.5$ mag. Hence we support the brighter luminosity scale for RR~Lyrae stars.  The range of metallicity needed to explain the whole extent of RRd band is [Fe/H]$=(-2,-1.3)$ for both Magellanic Clouds. While in SMC the objects appear to be uniformly distributed in this range, in LMC we see a strong concentration in the range $(-1.7,-1.3)$.  We have discussed uncertainties in our inference on metallicities following from the uncertainty of calculated values of period ratios. Its primary sources are treatment of convection and nonlinear effects as well as ambiguity of the heavy element abundance. All these effects contribute to uncertainty of [Fe/H] on the level of $0.5$. All of them should and can be reduced.  The observed width of the RRd band in the Petersen diagram may be explained by the spread in masses at given metallicity. Petersen diagrams provide a stringent constraint on RRd temperatures. As expected, these temperatures correspond to the mid of the RR~Lyrae range. In detail however, there is a difference between our RRd path in the HR diagram and that found by nonlinear modeling. Our RRd strip is significantly more inclined than the blue and red edges of the instability strip. The path from nonlinear simulations has inclination more than 90 degrees.  \\Acknow{We thank OGLE~Team for the unpublished data on RRd stars from SMC and G{\\'e}za~Kov{\\'a}cs for discussions. We also acknowledge useful comments of Giuseppe Bono. This work was supported in part by the Polish State Committee for Scientific Research grant 2-P03D-14. One of us (S.C.) was supported by MURST under the project ``Stellar evolution''.}"
    },
    "0011/astro-ph0011324_arXiv.txt": {
        "abstract": "The interactions between jets of high-energy nuclei and nuclei of the surrounding medium are studied. Such interactions could be initiated by jets from active galactic nuclei interacting with surrounding cool clouds. The resulting nuclear interactions are found to produce copious amounts of $^2$H and $^3$He from the $^4$He nuclei. These results suggest that jets of particles from quasars could have produced anomalously high abundances in surrounding clouds of some of the nuclides usually thought to characterize big bang nucleosynthesis, specifically, the $^2$H seen in absorption spectra. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011112_arXiv.txt": {
        "abstract": "We describe a near-infrared imaging survey covering $\\sim$ 1 deg$^2$ of the Chamaeleon I dark cloud.  The survey is complete for $K$ $<$ 15.0, $H$ $<$ 16.0, and $J$ $<$ 16.5, roughly two magnitudes more sensitive than previous large scale surveys. We use the large number of background stars detected to derive an accurate near-infrared extinction law for the cloud and select new candidate members with near-infrared color excesses.  We list $\\sim$ 100 candidates of the cloud with $K$ $\\ge$ 12.0, based on their positions in the $J-H$, $H-K$ color-color diagram.  These new stars have low luminosities ($K$ $\\sim$ 12 -- 16, $H-K$ $\\gtrsim$ 0.5 -- 1.5) and may have masses close to or even below the hydrogen burning limit. ",
        "introduction": "The Chamaeleon I dark cloud ($\\alpha$ $\\sim$ 11$^{\\rm h}$, $\\delta$ $\\sim$ $-$77$^{\\rm o}$) is an active stellar nursery with 150 or more known young stars \\citep{sch91,gast92,law96,cam98,com99}. Located at a distance of $\\sim$ 160 pc \\citep{whi97}, it has a relatively small angular size \\citep[$\\sim$ 3 deg$^{\\rm 2}$]{bou98}. Because of its proximity and moderate extension on the sky the Chamaeleon I star-forming region is an attractive place to attempt to measure the initial mass function (IMF), particularly at low masses (i.e., for masses \\hbox{$\\le$ 0.1 M$_{\\odot}$).}  Here we report results of a $JHK$ imaging survey of the Cha I dark cloud, which covers an area of $\\sim$ 1 deg${\\rm ^2}$ and is complete for $K$ $<$ 15.0, $H$ $<$ 16.0, and $J$ $<$ 16.5.  Our survey complements the $IJK_{\\rm s}$ observations of DENIS (Deep Near Infrared Southern Sky Survey, ESO, 1m tel, \\citet{epc97}) described by \\citet{cam98}.  The DENIS survey covers a larger area, 1.5$\\rm{^o}$ $\\times$ 3$\\rm{^o}$, of the cloud; but, with a sensitivity of $K_s$ $<$ 13.5, it is less sensitive to the lowest mass cloud members. \\citet{cam98} proposed 54 new candidate young stars, based on their locations in the $I-J$/$J-K$ diagram. In addition our survey area comprises two areas, in the the northern region, recently observed by \\citet{per99} and \\citet{oas99}. The data reported by these authors are $\\sim$ 2 magnitudes deeper than ours but cover modest areas on the cloud ($\\sim$ 3$'$ $\\times$ 3$'$, \\citet{per99}; $\\sim$ 6$'$ $\\times$ 6$'$, \\citet{oas99}) centered close to the high-velocity C$^{18}$O bipolar outflow previously detected by \\citet{mat89}. \\citet{per99} identified a new Class I source of the cloud. \\citet{oas99} proposed 9 new very low luminosity members of the cloud.  Our survey recovers 40 of these DENIS candidates (lying on common areas to both surveys) and the stars with $K$ $<$ 15.5 detected by \\citet{per99} and \\citet{oas99}. In addition the present observation yields $\\sim$ 100 new candidates with $K$ $\\ge$ 12 and near-infrared excess emission in the $J-H$,$H-K$ color-color diagram. The low luminosities of these candidates suggests they may be very low mass cloud members, with masses close to or below the hydrogen burning limit.  We describe our observations, data reduction, and analysis in \\S2. The astrometry and photometry were obtained for $\\sim$ 11,090 sources at $JHK$ detected in our survey region. In \\S 3 we derive a reliable extinction curve for background stars and use this curve to propose $\\sim$ 100 new potential young stellar objects based on their location in the $J-H$/$H-K$ diagram in \\S 4.  We conclude with a brief summary in \\S 5. ",
        "conclusions": "We have used the large number of background stars detected in our survey region to derived a reliable near-infrared extinction law for the Chamaeleon I cloud. Our analysis yields $E_{J-H}/E_{H-K}$ $=$ 1.80 $\\pm$ 0.03 for the Restricted Sample (i.e., eliminating previously known pre-main sequences stars).  This result differs with the \\citet{ke98a}'s determination for the $\\rho$ Oph cloud ($E_{J-H}/E_{H-K}$ $=$ 1.57 $\\pm$ 0.03) at more than 3$\\sigma$ level. Based on preliminary results for other two cloud (Taurus and IC 348), in addition to their determination for $\\rho$ Oph, \\citet{ke98a} suggest that a real variation in the near-infrared reddening law occurs from region to region. Our results for Cha I provide additional support for this variation.  The ratio of total to selective extinction, R$_V$, seems also to change for different clouds (see \\citet{ke98a} and the references therein).  In addition, R$_V$ varies across the Cha I dark cloud (from R$_V$ $\\sim$ 3 to R$_V$ $\\sim$ 5.5; \\citet{whi97}, see also \\citet{hyl82}).  For the $\\rho$ Oph, \\citet{vrb93} obtained a mean value of R$_V$ $\\sim$ 4, although individual determinations vary within a range similar to that found for the Cha I dark cloud. \\citet{hay99} recently found that the A$_V$ -- $N(CO)$ relation also varies among different star-forming clouds and suggest that A$_V$/$N(CO)$ may also change across the Chamaeleon I dark cloud. These results indicate different dust-grain properties, in particular grain-sizes, for different nearby star-forming regions and caution against using the `` universal '' reddening law to derive extinction estimates for young stars in star-forming regions.  Previous near-infrared observations of the cloud focussed on IRAS-selected sources \\citep{ass90,whi91,pru91} or already known and candidate young stellar objects detected at optical and X-ray wavelengths \\citep{gast92,har93,law96,com99}. These investigations were sensitive to objects with K $\\sim$ 11 and limited to small areas around the target objects. Present imaging near-infrared surveys of the cloud have overcome these limitations (both in sensitivity and spatial coverage) and greatly increased the number of potential members of the cloud. \\citep{per99} detected one new Class I source of the cloud with K $\\sim$ 13. \\citet{oas99} proposed 9 new candidates with K $=$ 13 -- 16.  \\citep{cam98} and this paper carry out large scale near-infrared survey on the cloud.  \\citep{cam98} proposed 54 new candidates members with K $\\lesssim$ 13. This paper presents an additional set of $\\sim$ 100 new stars with K $\\sim$ 12 -- 16. In addition to these candidates selected on basis of their near-infrared color excesses, \\citet{per00} have recently proposed a list of 74 new low mass candidates detected by ISOCAM and selected on basis of their color excesses at the mid-infrared wavelengths.  These objects are among the lowest luminosity and, presumably, lowest mass members of the cloud. Some of them are near or, probably, below the H burning limit. A spectroscopic follow up of the optically visible candidates should provide additional indications of their pre-main sequence nature. Optical and/or near-infrared spectral types of our proposed candidates, in combination with recently developed pre-main sequence evolutionary tracks, will allow us to determine masses and ages. Reliable mass determinations are required to calculate the IMF, particular towards the sub-solar and probably into the sub-stellar mass regime. Ages determination will help to reconstruct the history of the star-formation process of the cloud."
    },
    "0011/astro-ph0011438_arXiv.txt": {
        "abstract": "{ We present the XMM--Newton first light image, taken in January 2000 with the EPIC pn camera during the instrument's commissioning phase, when XMM--Newton was pointing towards the Large Magellanic Cloud (\\object{LMC}). The field is rich in different kinds of X--ray sources: point sources, supernova remnants (SNRs) and diffuse X--ray emission from \\object{LMC} interstellar gas. The observations are of unprecedented sensitivity, reaching a few 10$\\sp{32}$ erg/s for point sources in the \\object{LMC}. We describe how these data sets were analysed and discuss some of the spectroscopic results. For the SNR \\object{N157B} the power law spectrum is clearly steeper than previously determined from ROSAT and ASCA data. The existence of a significant thermal component is evident and suggests that \\object{N157B} is not a Crab--like but a composite SNR. Most puzzling is the spectrum of the \\object{LMC} hot interstellar medium, which indicates a significant overabundance of Ne and Mg of a few times solar. ",
        "introduction": "On December 10 1999 the XMM--Newton spacecraft was placed in a 48 hour Earth orbit by the first commercial ARIANE V launcher. XMM--Newton, or X--ray Multi--Mirror Mission (Jansen et al.\\ \\cite{J}), is the second cornerstone of the Horizon 2000 science program of the European Space Agency ESA.  The EPIC pn camera (Str\\\"uder et al.\\ \\cite{St}) was successfully commissioned during the period of mid January through mid March (Briel et al.\\ \\cite{Br}). On January 19 2000, shortly after the switch--on of the EPIC pn camera, XMM--Newton received first light on the EPIC pn CCD chip. We selected as first light target a region in the \\object{LMC} about 10 arcmin southwest of the \\object{30 Doradus~A} complex. The field has been known to be very rich in X--ray sources from Einstein (Long et al.\\ \\cite{L}) and ROSAT observations (Tr\\\"umper et al.\\ \\cite{Tr}). There are at least three SNRs (one of those known to contain a pulsar), one OB association and two R associations as well as the \\object{SN 1987A}. This XMM--Newton first light has already stimulated detailed studies of some of the objects. Here we report some results of the observation with four observation sequences of the same field added totaling an exposure time of 106~ks. The observations are of unprecedented sensitivity, reaching a few 10$\\sp{32}$ erg/s for point sources. We present and discuss the spectroscopic results over the broadest X--ray bandpass of 0.3\\,--\\,12~keV ever covered by a single imaging X-ray instrument for the entire field of 27$\\times$27 arcmin$^2$. In section 2 we describe details of the data analysis techniques which may be useful for other observations as well. In section 3 we discuss some of the scientific results. ",
        "conclusions": "As Fig.\\,\\ref{xmmlmc} shows, there is a wealth of different objects, point--like and extended up to 130~pc in size. The X--ray colours show large spectral differences from object to object. Based on this image, we accumulated photons from selected regions and performed spectral investigations of various sources, to utilize the full resolving power of the EPIC pn camera. We did this separately for the three data sets which were taken with the same observational setup (filter and pointing direction), because they require different detector response matrices. The best fit was obtained by modeling each of the three count rate spectra with the same candidate source spectrum but with the individual normalization factor left as a free parameter (for details see Haberl et al.\\ \\cite{Hab}). Errors were determined at $\\Delta\\chi^2=+2.3$ with all parameters free.  \\subsection{\\object{30 Dor C}}  The largest object in the field is \\object{30 Dor C}, already known from radio observations and suggested by Mills et al.\\ (\\cite{Mi}) to be an SNR. In X--rays it appears as a complete ring of $\\sim6$~arcmin diameter. The size of the shell and its colour make this remnant the first of its kind outside our Galaxy. In a separate paper (Aschenbach et al.\\ \\cite{A1}) this unique object is investigated in detail. Close to the geometric center we detect a previously unknown point source, displayed in yellow, which emits mainly in the 0.3\\,--\\,1.0~keV band. If this source is in the \\object{LMC}, then its luminosity is of the order of $10^{33}\\mbox{ erg s}^{-1}$, and it might be the compact remnant of \\object{30 Dor C}, either a neutron star or a black hole.  \\subsection{\\object{SN 1987A}}  Approximately 2~arcmin to the south-west of the shell border is the well--known \\object{SN 1987A}. The XMM--Newton observations are very important with respect to both, the light curve and the X--ray spectrum, refining current models of \\object{SN 1987A}. Both aspects are presented and discussed in an accompanying paper by Aschenbach et al.\\ (\\cite{A2}).  \\subsection{Active Galactic Nuclei behind the \\object{LMC}}  A number of point sources have been discovered, displayed as blue dots in Fig.\\,\\ref{xmmlmc}, the brightest of which is seen close to the image center. These sources have power law spectra and show strong low energy absorption. Both, the spectra and the fact that the central source is associated with a radio point source, indicate that we see background sources, probably AGN, Nuclei, through and far beyond the \\object{LMC}. A detailed analysis of these new X--ray sources is presented in an accompanying paper by Haberl et al.\\ (\\cite{Hab}).  \\subsection{\\object{N157B}}  By far the brightest object in the field is the SNR \\object{N157B}, located halfway from the image center to the edge in north--easterly direction. After the discovery of the 16~ms pulsar (Marshall et al. \\cite{Ma}, Wang \\& Gotthelf \\cite{Wa2}) \\object{N157B} has been classified as a Crab--like remnant with a central pulsar and a surrounding synchrotron nebula. In a preceeding paper Wang \\& Gotthelf (\\cite{Wa1}) showed that the spectrum is dominated by a power law with an indication of some thermal emission. The XMM--Newton spectrum of \\object{N157B} (Fig.\\,\\ref{n157b}) demonstrates convincingly for the first time the presence of emission lines at low energies, proving the presence of a thermal component. A model comprising two thermal components and a power law gives an acceptable fit to the data. Formally, ionization equilibrium models like Raymond--Smith (RS; Raymond \\& Smith \\cite{R}) or non equilibrium ionization models (NEI; Hamilton et al.\\ \\cite{Ham}) of XSPEC used for the thermal components at low energies cannot be discriminated.  Fig.\\,\\ref{n157b} shows the best fit two temperature RS models with a power law model added. The temperatures are 0.37$^{+0.11}_{-0.05}$ keV and 0.12$\\pm0.01$ keV, and the photon index of the power law is 2.83$\\pm$0.03. With an assumed hydrogen equivalent photoelectric absorption of $6\\cdot10^{20}$ H--atoms cm$^{-2}$ for the galactic foreground, we derive an additional H--equivalent column density of (1.94$\\pm$0.04)$\\cdot10^{22}$ H--atoms cm$^{-2}$ for the absorption in the \\object{LMC}, taking elemental abundances of 0.5 solar into account. To fit the Mg emission line at 1.3~keV clearly seen in the spectra, an overabundance of 1.3 solar is required (for the cooler component). The NEI model yields similar values for the absorption, the Mg abundance and the power law slope, but different temperatures (0.24 and 0.10 keV).  The existence of the additional two--temperature thermal component suggests that \\object{N157B} is not a Crab--like remnant but a composite remnant with an X-ray bright pulsar and synchrotron nebula but relatively faint thermal shell. The analysis of the thermal component, to be done in a separate paper, will be used to check for instance the consistency of the age derived from a Sedov--type approach with the pulsar characteristic spin-down age of $P/(2\\,\\dot P)\\sim5000\\mbox{ ys}$ (Wang \\& Gotthelf \\cite{Wa2}).  \\begin{figure}[tbp] \\resizebox{\\hsize}{!}{\\psfig{file=XMM54_f6.eps,angle=-90,clip=}} \\caption[ ]{EPIC pn X--ray spectra of N157B for the three data sets with the same observational setup. The histograms represent the best fit model simultaneously applied to the spectra.} \\label{n157b} \\end{figure}  The power law component in the spectrum of \\object{N157B} which has been discovered by Wang \\& Gotthelf (\\cite{Wa1}) using ROSAT PSPC and ASCA SIS data gives a photon index of 2.3$\\pm$0.1. The EPIC pn best fit value for the index is 2.83$\\pm$0.03. \\object{N157B} has been compared with the \\object{Crab}. Both the \\object{Crab pulsar} and the \\object{N157B 16~ms pulsar} (Marshall et al. \\cite{Ma}) show a flat power law spectrum with a photon index of 1.5 and 1.6$\\pm$0.3, respectively. The spectrum of the \\object{Crab nebula} is steeper by 0.5, generally believed to be due to synchrotron losses. If the spectrum of \\object{N157B} would have a similar difference of $\\sim0.5$ between the pulsar and the nebula spectrum as suggested by the analysis of Wang \\& Gotthelf, the analogy with the \\object{Crab} would be convincing. If on the other hand the EPIC pn power law slope is a better representation of the spectrum, the analogy with the \\object{Crab} is cleary violated and other acceleration and loss processes have to be considered. We note that the analysis of power law dominated spectra like in AGN show just minor discrepancies in the spectral index between ASCA and EPIC pn. Because of the steep spectrum of the nebula the physical conditions, either the electron injection spectrum, the electron transport speed or the magnetic field distribution, are likely to be different from those prevailing in the \\object{Crab}.  \\subsection{The \\object{Honeycomb nebula}} In Fig.\\,\\ref{xmmlmc}, X--ray emission from the \\object{Honeycomb nebula} is seen 1~arcmin south--east of \\object{SN 1987A}. The emission of the bright orange patch can be modeled by a thermal plasma with a temperature of 0.19$\\pm$0.02~keV, superimposed on a power law of index 3.6$\\pm$0.6. Increased Ne and Mg abundances ($\\sim$2 times solar) are indicated. The hydrogen equivalent LMC absorption to this SNR is (2.2$\\pm$0.9)$\\cdot10^{21}$ H--atoms cm$^{-2}$. Fig.\\,\\ref{honey} shows the three EPIC pn spectra together with the best fit model.  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\psfig{file=XMM54_f7.eps,angle=-90,clip=}} \\caption[]{Same as Fig.\\,\\ref{n157b} for the \\object{Honeycomb nebula}.} \\label{honey} \\end{figure}  \\begin{figure}[] \\resizebox{\\hsize}{!}{\\psfig{file=XMM54_f8.eps,angle=-90,clip=}} \\caption[ ]{EPIC pn X--ray spectrum of the faint, diffuse emission in the north of Fig.\\,\\ref{xmmlmc}, accumulated from the last two observations with thin filter.} \\label{s00} \\end{figure}  \\subsection{Interstellar medium in the \\object{LMC}}  The north--western part of Fig.\\,\\ref{xmmlmc} appears remarkably dark and empty. But when we extract a spectrum from a 44~arcmin$^2$ region north of \\object{30 Dor C} where no point sources and no out--of--time events from \\object{30 Dor C} affect the data, we clearly see thermal emission. The temperature appears to differ from place to place over a range from $\\sim$0.1 keV and $\\sim$0.3 keV, as was already concluded from ROSAT PSPC spectra (Sasaki et al.\\ \\cite{Sa}). But for the first time we detect emission lines of oxygen, neon, magnesium and silicon (Fig.\\,\\ref{s00}). To fit the spectrum, we subtracted the background in a preliminary way, using source free areas from other EPIC pn observations. The resulting spectrum is well represented by a 0.3$^{+0.05}_{-0.2}$ keV RS type model with overabundances in O, Ne, Mg and Si of 1.6$\\pm$0.7, 5.4$\\pm$1.5, 6.6$\\pm$2.5 and 8.9$\\pm$7.4 solar, respectively, which is surprising in view of the generally low metallicity of the LMC."
    },
    "0011/astro-ph0011262_arXiv.txt": {
        "abstract": "The new results of the evolution of the synchrotron peak for Mrk 421 are mostly likely related to the particle acceleration process. In order to account for the above results, we present a model of blazar variability during the flare in which the emission comes from accelerating electrons. A diffusion advection equation of the electron energy distribution is derived to calculate the spectrum and light curve of synchrotron radiation. We present that the observed shifts of the synchrotron peak moving to higher energies during the flare are caused by shock acceleration. The observed relation between changes in the fluxes at specified frequency ranges and shifts of the peak position is fitted to constrain the physical parameters of the dissipation region. ",
        "introduction": "Among active galactic nuclei, blazars are characterized by a high and variable degree of polarization and a flux variability often occurring on very short time scales. The observed continuum is dominated by nonthermal emission, where the emission deriving from synchrotron and self-Compton is enhanced by relativistic beaming (Blandford \\& Rees 1978). Some blazars exhibit not only intensity variation extending from radio waves to gamma rays, but also spectral variation as a function of the flux level in many frequency bands.  An interpretation of the spectral variations can be envisaged in an inhomogeneous jet scenario of the kind proposed by Ghisellini et al. (1985) and Celotti et al. (1991). In this model the structure of the jet is considered, the magnetic field and the relativistic electron density are assumed to decrease with distance, and the maximum synchrotron frequency is assumed to decrease with distance, accounting for a variability time shorter at larger synchrotron frequencies. The predicted variability pattern is caused by a disturbance (e.g. shock) traveling down the jet. This perturbation is assumed to produce only a fixed enhancement in magnetic field and particle density and to unchange the shape of the electron distribution. The overall spectrum is the superposition of the located spectra emitted by each slice of the jet. As it has been applied to BL Lac objects, the model does not consider the evolution of the relativistic electron distribution affected by radiation losses, new injection of relativistic electrons and shock waves. However, simultaneous observations of blazars over a wide spectral range suggest that significant radiation losses and injection of relativistic electrons occur in flares. For example, the variability of PKS2155-304 at X-ray, EUV and UV bands is decreased in amplitude and with significant delays approximately satisfying a $\\nu ^{-\\frac 12}$ relation (Urry et al. 1997). In particular, the radiation loss time for electrons that emit optical and X-ray emission in blazars is probably far less than the light (or the shock) crossing time of the emitting region. Consequently, if the electron distributions responsible for high energy emission are kept by the shock acceleration processes, then radiation losses and the injection of electrons occur in the vicinity of the shock during the flare and must be considered in the theory. Many models to reproduce the spectral variability have been developed ( Mastichiadis \\& Kirk 1997; Georganopoulous \\& Marscher 1998; Kirk, Rieger \\& Mastichiadis 1998; Wang et al. 1999; Ghiaberge \\& Ghisellini 1999; Li \\& Kusunose 2000). The observed spectrum and time delays between the light curves at fluxes at different frequencies are believed to be produced by the electron distribution at different stages of evolution after episodic electron injection phases. However these models can not account for a synchrotron peak drifting to higher energies at X-ray wavelengths during the rising phase of the flare in Mrk 421 (Fossati et al. 2000). The above hard lag is most likely related to the particle acceleration process. The role of particle acceleration by shock waves has been considered by Kirk et al. (1998). In this model, electron acceleration and radiation zones are separated. Electrons are continuously accelerated at the shock front, and subsequently drift away from it in the downstream and emit most of the radiation. The emission from the acceleration zone is ignored.  Through shocks can quickly accelerate particles to very high energies, this requires the existence of some scattering agent to force repeated passage of the particles across the shock. The most likely agent for scattering is plasma turbulence or plasma waves. However, the plasma turbulence needed for the scattering can not only accelerate particles stochastically (second order Fermi acceleration), but can cause particle diffusion in the downstream zone quickly. If the dissipation of the bulk energy of blazar jet into particles includes shocks and plasma turbulence, it seems difficult to distinguish acceleration zone and emission region. The emitting particles in downstream could be re-accelerated by shock waves under the scattering of plasma turbulence. Therefore we try to consider the synchrotron emission from a single dissipation region which includes shock and turbulence processes and radiation losses.  In Section II we present a stochastic differential equation to describe a particle acceleration and cooling processes, together with the assumptions made. Then we derive a diffusion-advection equation of particle energy distribution which simulates the temporal evolution of the particle radiation occurring when shock waves and plasma turbulence are produced in a dissipation region of a jet. In section III we apply our model to Mrk 421 to explain the hard lag observed in the X-ray light curves. Finally, we draw our conclusions in section IV. ",
        "conclusions": "The recent BeppoSAX observations of Mrk 421 (Fossati et al. 2000) provide important information to understand particle acceleration processes. Within a single emission region scenario for blazar jets, we have studied the time dependent behavior of the particle distribution affected by the particle acceleration and radiation losses, and calculated the form of light curves and spectra at different times. We have presented that the observed shifts of the synchrotron peak moving to higher energy during the flare are caused by shock acceleration. The accelerating particles follow a narrow Guassian distribution with a central energy quickly evolving with time to the highest energy when the shock acceleration dominates the turbulent acceleration. The highest energy is limited by radiation losses. Our results are important for the observed fast dissipation region of blazar jets where the light crossing time of the region is shorter than the particle acceleration and cooling time scales. The observed fast variability indicates the particle acceleration and cooling processes.  The observed relation between changes in the fluxes at specified frequency ranges and shifts of the peak position during the flare can estimate the magnetic field and the number density of accelerated particles in the dissipation region.  An energy dependence of the shape of the light curve observed in Mrk 421 during the flare (Fossati et al. 2000) can be connected to the particle acceleration and cooling time scales. During the acceleration, the observed light curves are expected to be symmetric at low energies where the acceleration time is similar to the cooling time and much longer than the light crossing time. The asymmetric light curves (faster rise) occur at higher energies where the acceleration time is shorter than the cooling time and is comparable to the light crossing time. During the post-acceleration, the light curve is symmetric where the rise and decay timescales are equivalent to cooling time scale.  With the steady state solution of the diffusion-advection equation of particle energy distribution, we have demonstrated that the turbulent acceleration leads the particle energy diffusion. The strength of turbulent acceleration determines the width of energy diffusion which modifies the particle distribution with respect to the pure shock acceleration case."
    },
    "0011/astro-ph0011040_arXiv.txt": {
        "abstract": "Motivated by hydrodynamic simulations, we discuss the X-ray appearance of radio galaxies embedded in the intracluster medium (ICM) of a galaxy cluster.  We distinguish three regimes.  In the early life of a powerful source, the entire radio cocoon is expanding supersonically and hence drives a strong shock into the ICM.  Eventually, the sides of the cocoon become subsonic and the ICM is disturbed by the sonic booms of the jet's working surface.  In both of these regimes, X-ray observations would find an X-ray shell.  In the strong shock regime, this shell will be hot and relatively thin.  However, in the weak shock (sonic-boom) regime, the shell will be approximately the same temperature as the undisturbed ICM.  If a cooling flow is present, the observed shell may even be cooler than the undisturbed ICM due to the lifting of cooler material into the shell from the inner (cooler) regions of the cluster. In the third and final regime, the cocoon has collapsed and no well-defined X-ray shell will be seen.  We discuss ways of estimating the power and age of the source once its regime of behavior has been determined. ",
        "introduction": "In the powerful Faranoff-Riley type II radio galaxies (FR-II; Fanaroff \\& Riley 1974), jets from a central AGN propagate at relativistic speeds before terminating in a series of shocks resulting from interaction with the surrounding material.  The spent jet material is then thought to inflate a broad cocoon which encases the jets and is seen as the observed radio lobes.  In very powerful or young sources, the cocoon is probably overpressured (Begelman \\& Cioffi 1989) with respect to the surrounding ambient material (either the interstellar medium [ISM] of the host galaxy, or the intracluster medium [ICM] of the host cluster).  The cocoon then undergoes supersonic pressure driven expansion into the ambient medium, sweeping the ambient medium into a shocked shell.  A contact discontinuity separates the relativistic cocoon material from the shocked ambient material.  X-ray observations provide one with a direct probe of the interaction of the ISM/ICM with the radio galaxy cocoon.  X-ray imaging observations with the high resolution imagers (HRIs) on the {\\it ROSAT} satellite discovered several systems in which the observable radio lobes are seen to reside in evacuated cavities within the ICM which are bounded by shells of enhanced emission (e.g. Cygnus-A: Carilli et al. 1994; Perseus-A: Bohringer et al. 1992; A~4059: Huang \\& Sarazin \\& 1998).  In order to further explore the shell/cavity in Cygnus-A, Clarke, Harris \\& Carilli (1997) performed 3D hydrodynamic simulations of a jet propagating into a cluster-like atmosphere.  They found that the X-ray cavities are readily identified with the radio cocoon (whose pressure driven expansion excludes the ICM), and that the X-ray shells can be identified with strong shocks in the ICM.  Kaiser \\& Alexander (1999) showed how analytic self-similar hydrodynamic models can be used to obtain similar results.  More recently, Rizza et al. (2000) performed 3-D hydrodynamic simulations in an attempt to model subtle X-ray structure in A~133, A~2626, and A2052.  In a complementary study, Heinz, Reynolds \\& Begelman (1998; hereafter HRB98) used a simple ``bubble'' model for the cocoon/shock system (Begelman \\& Cioffi 1989; Reynolds \\& Begelman 1997) to predict the X-ray appearance of radio galaxies as a function of their age and power.  A central assumption used in the HRB98 model is that the presence of an X-ray shell bounding the {\\it sides} of a radio cocoon indicates a lateral shock front, which in turn implies that the sides of the cocoon are still expanding at supersonic speeds.  This assumption is based upon the hydrodynamical fact that, once the sideways expansion of the cocoon becomes subsonic, the contact discontinuity separating the cocoon and the ICM will become susceptible to the Rayleigh-Taylor instability on timescales shorter than the expansion timescale of the cocoon.  One might then expect the cocoon to collapse, with the corresponding disappearance of the X-ray cavities.  In this letter, we follow up the study of HRB98 by using moderate-to-high resolution 2-D hydrodynamic simulations to study the interaction of radio-galaxies with the ICM.  We find that powerful radio galaxies in which the entire cocoon is supersonic do indeed produce a well defined shell/cavity structure, although backflows within the shocked ICM shell increase its thickness substantially over that predicted by the bubble model (also see Kaiser \\& Alexander 1999). However, we also find a shell/cavity structure in sources in which the sides of the cocoon are expanding at subsonic speeds (or even not expanding at all).  In this case, the cocoon has not had time to collapse, despite the fact that Rayleigh-Taylor instabilities are gradually destroying it.  The X-ray shell corresponds to emission from a weak shock or compression wave (i.e. a sonic boom) which is driven into the ICM by the supersonic motion of the cocoon head.  We find that ICM material can be lifted into this emitting shell without passing through a strong shock, which has important implications for the temperature structure of such X-ray shells and entropy injection into the ICM.  This study is particularly timely given its relevance to {\\it Chandra} observations of radio galaxies embedded within galaxy clusters.  Section~2 briefly discusses the configuration and results of our hydrodynamic simulations.  In Section~3, we compute emissivity maps and discuss the nature of X-ray shells and cavities.  Section~4 describes how one may use X-ray data to estimate the power and age of a given radio galaxy.  We summarize our conclusions in Section~5. ",
        "conclusions": "Guided by 2-D hydrodynamic simulations, we distinguish three regimes of radio galaxy evolution which have importance for understanding the structure of X-ray cavities and shells.  In the early life of a powerful radio galaxy, the entire cocoon is expanding supersonically and a strong shock is driven into the ICM in all directions.  One would see a relatively narrow and hot X-ray shell surrounding the radio lobes.  However, eventually, the lateral expansion velocity of the contact discontinuity becomes subsonic.  There can be a substantial phase of the source's life in which the cocoon is subsonic but has yet to collapse and is surrounded by a rather thick shell of disturbed ICM which has suffered a weak shock or compression wave. X-ray observations of such a system would reveal a rather thick shell bounding a cavity.  The temperature of this shell would be very similar to the temperature of the undisturbed ICM.  If a cooling flow is present, this shell would contain cooler material that has been lifted from further down in the cluster potential.  Even if the ICM is isothermal, regions of the disturbed shell could be slightly cooler due to the action of adiabatic expansion.  We also note that a backflow exists in the disturbed ICM shell which significantly depletes the shocked shell near the head of the jet and renders it difficult to observe.  After several ICM dynamical timescales, the cocoon will collapse and the apparent X-ray shell will vanish.  Finally, at very late times, after the radio galaxy activity has ceased, the relic lobes may appear as `naked' cavities within the ICM emission (i.e. X-ray cavities that are not bounded by shells).  Such structures have been explored in recent work by Churazov et al. (2000)."
    },
    "0011/astro-ph0011089_arXiv.txt": {
        "abstract": "The RXTE observations of GRS 1915+105 have given a new impulse to the study of spectral and timing properties of the X-ray emission of black hole candidates. At variance with any other known source, GRS~1915+105 shows dramatic changes on time scales as short as a tenth of a second. These changes are associated to marked spectral changes which have been interpreted as changes of the observable inner radius of the accretion disk. I review the existing results and discuss the current evidence for such disk oscillations. Independently of the precise theoretical models, the detailed study of these fast variations can provide an extremely valuable insight on the accretion processes onto black holes. Making use of these results, I compare the properties of GRS~1915+105 with those of other black hole candidates. ",
        "introduction": "In the recent years, a classification scheme for the X-ray emission of black hole candidates has emerged. From the timing and spectral properties, four separate states (plus the quiescent state for transient systems) have been identified, probably in dependence of the accretion rate level (see van der Klis 1995). Unfortunately, all known persistent sources, Cyg X-1, GX~339-4, LMC~X-1 and LMC~X-3, show state transitions very rarely, making the accumulation of extensive datasets on different states difficult. Transient systems are more promising, but their transient nature limits the number of states and state transitions that can be observed from one system, and the comparison between different systems is always problematic.  The spectral properties of black hole candidates are usually characterized in terms of a double-component model. The first component is thermal and (when present) contributes mostly to the flux below 10 keV. It is commonly interpreted as emission from an optically thick accretion disk (Mitsuda et al. 1984). The second component is much harder and extends above 10 keV. In its simple form, it can be approximated as a power law, although more sophisticated models are often required (see e.g. Frontera et al. 2000). An attractive feature of the optically thick disk model, the so-called {\\it disk-blackbody}, is that one of its parameters is the value of the inner radius of the accretion disk. Therefore, the application of this model can in principle yield the measurement of this fundamental parameter, although a precise value depends on the distance and the inclination of the system. This model became rather popular when a number of black hole systems showed a rather constant inner disk radius, around 20-30 km, despite large changes in accretion rate (see Tanaka \\& Lewin 1995). Although this model is simplified and leads to an underestimate of the values of the radii, it can detect large variations in the inner radius of the disk (see Merloni, Fabian \\& Ross 2000). ",
        "conclusions": "GRS~1915+105 is the only source up to now which can provide a large number of spectral transitions, enabling us to study different states of one source without having to wait for years. The most important of these transitions do not involve the canonical states of black hole candidates, but are associated to an instability of the innermost region of an optically thick accretion disk, which causes its inner parts to be evacuated and refilled on the local viscous time scale. Other variations observed in this source can be linked to those observed in other sources, strengthening the connection between this unique system and more conventional ones. On the X-ray side, it is now important to examine the spectral/timing behavior of GRS 1915+105 in more detail (see Migliari, Vignarca \\& Belloni, this volume), in order to provide a more complete phenomenological picture for further theoretical work. Modeling of the X-ray properties of GRS~1915+105 is at the moment the most promising way to make significant progress in the understanding of accretion onto black holes.   \\theendnotes"
    },
    "0011/astro-ph0011276_arXiv.txt": {
        "abstract": "Anomalously high K-band surface brightness fluctuations (SBF) have been reported in NGC 4489 by Pahre \\& Mould (1994), Jensen et al. (1996) and Jensen et al.  (1998).  However, these conclusions were uncertain because of relatively low signal-to-noise data.  New high signal-to-noise data for NGC 4489 have been obtained at the NOAO/KPNO 2.1m and the ESO/La Silla 3.5m NTT telescopes.  Adopting the I-band SBF distance modulus determined by Tonry et al.  (1990) and the $\\overline M_I$ versus (V--I) calibration of Tonry et al.  (2000), a value of $\\overline M_K =-6.18\\pm0.09 $ mag was derived.  Relative to the average empirical $\\overline M_K$ derived for giant ellipticals by Jensen et al.  (1998) ($\\overline M_K=-5.61\\pm0.12$), the detection of an anomalous K-band SBF in NGC 4489 is confirmed at the two sigma level.  Such anomalous fluctuations could be caused by an extended giant branch, consisting of either intermediate-age AGB stars above the tip of the first-ascent giant branch or high-metallicity first-ascent giants, or by lack of a full understanding of the K--band SBF calibration. This result raises questions about the accuracy of K-band SBF distance measurements for elliptical galaxies with unknown stellar composition and underscores the need for $\\overline M_K$ measurements over a larger range of color and luminosity. ",
        "introduction": "The Poissonian statistics of stellar populations in each pixel of a galaxy image produces fluctuations. While the mean flux per pixel does not depend on distance, the variance of these fluctuations is inversely proportional to the square of the galaxy distance. When properly calibrated this variance can then be used as a distance indicator.  Surface Brightness Fluctuations (SBF) are defined as the variance of the flux of the fluctuations normalized to the mean flux of the galaxy in each pixel \\cite{ts88}. The stars that contribute the most to the SBF are then the brightest stars in the intrinsic luminosity function, typically red giants in old stellar populations (a recent review on the method has been given by Blakeslee et al. 1999).  \\noindent I--band SBF have being used successfully to measure elliptical and S0 distances up to 4000 km s$^{-1}$ from ground based telescopes and 7000 km s$^{-1}$ on the Hubble Space Telescope (HST) \\cite{sod95,sod96,aj97,ton97,tho97,lau98,pah99,bla99,ton00}.  Recently SBF measurements have been attempted in the infrared K--band for several reasons \\cite{lup93,pah94,jen96,jen98,jen99}.  \\noindent First of all the contrast between brightest red giants and the underlying stellar populations is more extreme in the K--band, producing larger SBF at any given distance. Second, at ground facilities, K--band seeing is intrinsically better than I--band seeing which also enhances the SBF contrast. Moreover the color contrast between K--band SBF and external sources (globular clusters and background galaxies) is larger than in the I--band.  Finally both empirical and theoretical studies suggest that absolute K--band SBF is almost constant with age for solar-metallicity stellar populations \\cite{pah94,jen96,jen98,jen99,mei99,liu00,mei00}.  \\noindent These arguments suggest that in the infrared it should be possible to determine SBF distances for galaxies at larger distances than at optical wavelengths \\cite{jen98}.  However, the K-band has one major disadvantage for ground-based studies - the background is significantly brighter making it much more difficult to detect and correct for globular cluster and background galaxy contamination and increasing the background shot noise.    \\begin{figure*}  \\centerline{\\psfig{file=4489f1.eps,width=10cm,height=8cm}}  \\caption{NGC 3379 external source luminosity function derived from the NTT data.  The continuous line shows the sum of the adopted globular cluster plus background galaxy luminosity functions.  The dashed line shows the globular cluster luminosity function, while the dotted line shows the background galaxy luminosity function.  In total, 38 sources were detected between $\\approx$ 15$\\arcsec$ and $\\approx$ 60$\\arcsec$.  } \\label{fig-3379gc}  \\end{figure*}    \\noindent Initial K-band SBF studies suggested that the $\\overline M_K$ varied very little from galaxy-to-galaxy \\cite{pah94,jen96}. However, these studies also argued that two Virgo ellipticals (NGC 4365 and NGC 4489) had significantly brighter than average $\\overline M_K$ values suggesting the presence of an extended giant branch.  This conclusion, however, was weakened by two problems: (1) no correction was made for globular cluster and background galaxy contamination; and (2) the SBF measurements were based on relatively low signal-to-noise data.  Jensen et al.  (1998) re-measured $\\overline M_K$ for NGC 4365, using I-band data to determine a correction for globular cluster contamination, and concluded that NGC 4365 was a member of the background W cloud and the $\\overline M_K$ for this galaxy was consistent with other Virgo ellipticals.  In short, NGC 4365 appears to have a normal giant branch.  No additional study of NGC 4489 has been published until now.  \\noindent In this study, new observations with significantly higher signal-to-noise and a new SBF analysis for NGC 4489 are presented.  The observations are described in $\\S$ 2, the SBF analysis in $\\S$ 3 .  The results of this analysis are presented in $\\S$ 4 and $\\S$ 5. ",
        "conclusions": "New, high signal-to-noise K-band SBF measurements have been obtained for NGC 3379 and NGC 4489.  In summary:  \\begin{itemize}  \\item K--band SBF $\\overline M_{K}$ values of $-5.32 \\pm 0.13$ mag for NGC 3379 and $-6.18 \\pm 0.09$ mag for NGC 4489 have been derived. Although the new apparent SBF magnitudes presented here are independent observations and are more accurate for NGC 4489 than previous determinations, the final K-band SBF absolute magnitudes rely on the accuracy of the respective distance moduli determined in the I-band.  \\item Keeping in mind the caveat above about absolute magnitudes, the measurement by Jensen et al.  (1998) and Pahre \\& Mould (1994) of large K--band SBF in NGC 4489 has been confirmed.  These fluctuations could be the signature of an extended giant branch, but this cannot be confirmed until K-band SBF magnitudes are determined for a larger sample of galaxies similar in color to NGC 4489.  \\end{itemize}  These results raise questions concerning the accuracy of K--band SBF distance measurements for elliptical galaxies.  SBF measurements have been extended to this band in the belief that one could determine with the same accuracy K--band SBF distances to larger distances than possible at optical wavelengths.  This is possible only if the contribution of external sources to the observed image fluctuations can be correctly estimated and if the K--band SBF absolute magnitude can be precisely calibrated.  The result of this paper concerns the issue of calibration.  The presently small sample of K--band measurements (around 20) would seem to suggest that $\\overline M_K$ varies very little from galaxy to galaxy.  However, our result -- a galaxy with fluctuations different than the current average, illustrates the importance of an extended K--band calibration over a larger range of luminosity and color.  Only larger samples will allow us to determine whether or not NGC 4489 (or for that matter M32) is an exceptional case, and to understand the real potential of K--band SBF distance determinations."
    },
    "0011/astro-ph0011106_arXiv.txt": {
        "abstract": "The lack of a high energy cutoff in the cosmic ray spectrum together with an apparently isotropic distribution of arrival directions for the highest energy events have strongly constrained most models proposed for the generation of these particles.  An overview of the theoretical proposals are presented along with their most general signatures.  Future experimental tests of the different proposals are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011330_arXiv.txt": {
        "abstract": "\\tighten  We have made a deep NICMOS imaging study of a sample of 5 $z \\sim 2$ -- 3 radio-quiet quasars with low absolute nuclear luminosities, and we have detected apparent host galaxies in all of these. Most of the hosts have luminosities approximately equal to present-day L$_*$, with a range from 0.2 L$_*$ to about 4 L$_*$. These host galaxies have magnitudes and sizes consistent with those of the Lyman break galaxies at similar redshifts and at similar rest wavelengths, but are about two magnitudes fainter than high-$z$ powerful radio galaxies. The hosts of our high-$z$ sample are comparable to or less luminous than the hosts of the low-$z$ radio-quiet quasars with similar nuclear absolute magnitudes. However, the high $z$ galaxies are more compact than the hosts of the low $z$ quasars, and probably have only 10 -- 20\\% of the stellar mass of their low-$z$ counterparts. In one host, we find a residual component that is not centered on the quasar nucleus, and several hosts have apparent companions within a projected distance of $\\sim$10 kpc, indications that these systems are possibly in some phase of a merger process. Application of the $M_{bulge}$/$M_{BH}$ relation found for present-day spheroids to the stellar masses implied for the high $z$ host galaxies would indicate that they contain black holes with masses around $10^{8}$ $M_\\odot$. Comparison to their nuclear magnitudes implies accretion rates that are near or at the Eddington limit. Although these high $z$ hosts already contain supermassive black holes, the galaxies will need to grow significantly to evolve into present-day $L_*$ galaxies. These results are basically consistent with theoretical predictions for the hierarchical buildup of the galaxy host and its relation to the central supermassive black hole. ",
        "introduction": "Radio-quiet quasars (RQQ) represent the great majority of the known active galactic nuclei (AGN) at high redshift ($z \\gtrsim 1$), yet information on the galaxies that host these AGN is currently extremely limited. At lower redshift, luminous RQQ are found in both spiral and elliptical hosts, with the most luminous radio-quiets appearing in the bulge-dominated systems. The host galaxy luminosities correspond to a few times a present-day $L_*$ galaxy (e.g. Bahcall et al. 1997; McLure et al. 1999). In addition, the luminosity of the host seems to correlate roughly with that of the nucleus (McLeod \\& Rieke 1995; Hooper et al. 1997; McLure et al. 1999; McLeod, Rieke, \\& Storrie-Lombardi 1999). Nearby, almost every bulge-dominated galaxy may contain a supermassive black hole candidate whose mass is roughly proportional to the luminosity of its bulge (Kormendy \\& Richstone 1996; Magorrian et al. 1998; van der Marel 1999). The black hole masses are even better correlated with the stellar velocity dispersions observed in the bulges (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a). These relationships imply a close correlation between the mass of the stellar bulges in these systems and the masses of their resident black holes.  The evolution in the star formation rate from an epoch of $z \\sim 2$ to the present day may be qualitatively similar to the huge drop in the number density of quasars from the peak at $z \\sim$ 2 -- 3 to the present epoch, which is roughly similar to the evolution in the comoving number density of radio sources (e.g, Dunlop 1998; Boyle \\& Terlevich 1997). All of these factors suggest a strong link between the formation and evolution of galaxies and those of the quasars and their hosts, and support the idea that an active quasar may be a short-lived episode in the life of every bulge-dominated galaxy.  Kauffmann \\& Haehnelt (2000) have investigated the evolution of the hosts of the radio-quiet quasars and the formation and fuelling of their central black holes using semi-analytic hierarchical clustering models of galaxy formation, and have made specific predictions regarding the evolution of this nuclear magnitude--host galaxy luminosity relation. The high-$z$ radio-quiet quasars are far less studied than the radio-loud host galaxy population, and few observations have been made at nuclear luminosities similar to those of the low-$z$ radio-quiet sample.  To study the evolution of the radio-quiet quasar hosts and their relationship to the other samples of high-$z$ galaxies known, we have made HST NICMOS observations of radio-quiet quasar hosts near the epoch of the peak quasar density ($z \\sim$2 -- 3). We have chosen a sample of quasars whose nuclei are faint enough to provide a good comparison sample to those of the well-studied low-$z$ quasars. Throughout this paper we use  $H_0$ = 50 km s$^{-1}$ Mpc$^{-1}$ and $\\Omega$  = 1, unless otherwise stated. ",
        "conclusions": "We have detected host galaxies around all five of our sample of $z \\sim 2$ -- 3 radio-quiet quasars. The detected hosts have absolute $V$ magnitude $M_V$ $\\sim$ $-$20.2 to $-$22.5, with absolute luminosities from less than present-day $L_*$ to 4 $L_*$. The quasar nuclei are of relatively faint absolute $B$ magnitudes, with $M_B$ ranging from $-$24 to $-$22. The scale sizes, where we can fit the residuals, yield half-light radii of $<$4 kpc.  \\subsection{Host galaxy magnitudes compared}  Observational resolution is obviously important in studying the hosts of quasars, particularly those at high redshift, since (at $z$ = 2) $\\sim8$ kpc of physical scale would be hidden under a typical $1\\arcsec$-diameter groundbased seeing disk. We will thus compare our work primarily to the results of other HST imaging; however, work has been done on high-$z$ quasars with ground-based resolution that provides tantalizing hints of possible correlations. While Lehnert et al. (1992) found bright hosts around 6 $z \\sim 2.5$ radio-loud quasars in ground-based near IR imaging, Lowenthal et al. (1995) detected no hosts around a comparison sample of RQQ at the same $z$ and nuclear magnitudes. This result implied that the host luminosity might correlate with radio-loudness.  In contrast, Aretxaga et al.\\ (1995, 1998) studied several very luminous $z \\sim 2$ RQQs and found hosts comparable in luminosity to those found in the RLQ sample of Lehnert et al. (1992), indicating perhaps a heavy dependence of host magnitude on the nuclear luminosity.  Our observations of the $z \\sim 2$ -- 3 RQQ sample benefit from the compact and stable NICMOS PSF. Indeed, our simple simulations with observations of other high-$z$ galaxies have indicated that for most intrinsic galaxy morphologies our estimates of the magnitudes and morphologies are not badly affected by the presence of the quasar nucleus and observing process.  We have quoted in Table \\ref{maghosttab} simple aperture magnitudes, with error estimates based on the variation in the subtraction and from the sky noise. In the last column, we also give a rough estimate of ``total'' magnitude, based on adding the systematic subtraction and aperture corrections discussed in section \\ref{sec:apert} to the 2\\arcsec\\ aperture magnitudes. It is these estimates of total magnitude that we will use in comparing our host galaxy properties to those of other samples of galaxies.  We plot in Fig. \\ref{mabsvfig}, for two different cosmologies, the absolute rest-frame $V$ total magnitudes of the 5 MZZ quasar hosts, extrapolating our $z \\sim 2.7$ quasar host $M_B$ values to $M_V$ assuming the hosts have a spectral index of $-$2 (i.e. flat in F$_\\lambda$), corresponding to a ($B$-$V$) of 0.7.  We also correct to rest-frame $M_V$ and plot on these figures the Lyman break galaxy magnitudes from the Dickinson (2000) sample, the powerful MRC/USS radio galaxies of the Pentericci et al. (2000) sample, and at low-$z$, the quasar host magnitudes from the Bahcall et al. (1997) and McLure et al. (1999) samples. In addition, Lacy et al. (2000) have made ground-based $K'$ imaging studies of a $z \\sim 2$ -- 3 sample of faint 7C-III radio galaxies for which we have both small and large aperture magnitudes, and we have included these as well.  We have used in all cases simple estimates of total magnitudes, applying corrections to the available aperture magnitudes of generally less than 0.3 mags. For the HDF LBG galaxies, we measured the magnitudes directly from the image provided by Dickinson et al. (2000). As these galaxies have similar compactness to our host galaxies, we used the same metric apertures used for our quasar hosts (i.e. 2\\arcsec\\ at $z = 1.8$ and $z = 2.7$, corresponding to $\\sim$ 16 kpc), and made similar aperture corrections to total magnitude as those used for the MZZ quasar hosts. For the Pentericci et al. (2000) bright radio galaxy sample, we have corrected their quoted 4\\arcsec\\ aperture magnitudes to total magnitudes assuming the hosts are ellipticals with scale sizes of 10 kpc (corrections of $\\sim$0.25 mags). For the 7C-III radio galaxy sample of Lacy et al. (2000), we correct the quoted 64 kpc aperture magnitudes to total, with the same assumptions (resultant corrections $\\sim$0.15 mags). The low-$z$ quasar host magnitudes of Bahcall et al (1997) and McLure et al. (1999) are discussed in more detail in section \\ref{sec:lowzsim}, but do not require correction to total magnitude. All magnitudes were observed at close to rest-frame $B$ or $V$ and therefore the $K$-corrections are minor.  The MZZ quasar host galaxy magnitudes at both $z$ ranges are much more consistent with those of the Lyman break galaxies at high-$z$ than with either sample of radio galaxies. At $z \\sim 1.8$, our mean $M_V$ is about $-$21.5, while the radio galaxies have mean $M_V$ $\\sim$ $-24.0$. At higher $z$, the difference is less, particularly between the less luminous 7C-III galaxies and our hosts, but the means still differ by almost 2 magnitudes. When compared to the Lyman break galaxies, the MZZ hosts span the same range of magnitudes over this redshift range.  Our derived magnitudes are similar to those of the low-$z$ quasars: slightly less luminous in the high $\\Omega_m$ cosmology, and comparable in the low $\\Omega_m$ cosmology. However, the stellar masses corresponding to these absolute $V$ magnitudes will be much less for the young, high-redshift galaxies than for the old stellar populations probably associated with the low-$z$ galaxies. We have placed models of the passive evolution of some simple stellar populations on the figures.  These models have been generated with the most recent of the Bruzual \\& Charlot population synthesis models (Bruzual \\& Charlot 2000), using a Salpeter IMF and an upper mass cutoff of 100 M$_\\odot$, with metallicity $Z = 0.02$ (Charlot 2000, priv. comm; further discussion of these models can be found in Bruzual \\& Charlot 1993; Liu, Charlot, and Graham, astro-ph/0004367). The dashed line represents a model with an instantaneous burst of star formation at $z = 5$, while the dot-dashed line represents a model with a 1 Gyr burst of star formation ending at $z = 3$. The evolutionary tracks are roughly normalized to the absolute $V$ magnitudes at $z = 1.8$ of (1) the MZZ quasar hosts and 2) the powerful radio galaxies. This corresponds to total masses in {\\it e.g.}, the instantaneous burst model, of (1)  1.1 $\\times$ 10$^{11}$ M$_\\odot$ ($\\Omega_m$=1) or 2.0 $\\times$ 10$^{11}$ M$_\\odot$ ($\\Omega_m$ = 0.3) and (2) 1.2 $\\times$ 10$^{12}$ M$_\\odot$ ($\\Omega_m$=1) or 1.9 $\\times$ 10$^{12}$ M$_\\odot$ ($\\Omega_m$ = 0.3). The total stellar masses for the 1 Gyr burst models are about 20 -- 30\\% smaller. The degree of fading is about 2.2 -- 2.5 magnitudes in $V$ from $z=2.7$ to $z \\sim 0$ and 1.5 -- 1.7 magnitudes from $z = 1.8$ to $z \\sim 0$, for the range of models and cosmologies considered, corresponding to (M/L) increases of $\\sim$ 10 from $z = 2.7$ to $z \\sim 0$. In other words, at $z = 0.2$, our mean quasar host would be about 2 magnitudes fainter than the mean $z = 0.2$ quasar host from the McLure et al. and Bahcall et al. samples if it evolved passively from a short burst of star formation at $z = 3$ -- 5. This implies that the mean MZZ quasar host may have about one sixth the stellar mass of the mean low-$z$ quasar host.  Thus we find that the hosts of $z \\sim 2$ -- 3 radio quiet quasars with faint nuclei ($M_B$ $\\sim$ $-$24 to $-$22) have magnitudes that are consistent with their being drawn from the population of Lyman break galaxies observed at similar redshifts.  Our results are consistent with other recent NICMOS results on some brighter, lensed quasars; Rix et al. (1999) find that their sample of lensed $z \\sim 2$ radio-quiet quasars (de-magnified $M_B$ $\\sim$ $-$24 to $-$28) had hosts with comparably faint magnitudes. In addition, first results from the large NICMOS quasar host survey of Kukula et al. (2000) have been reported, giving similar or slightly brighter host magnitudes for a sample of 5 $z \\sim 2$ RQQs with nuclear $M_B \\sim$ $-$24.  These magnitudes are not consistent, however, with those of the sample of radio galaxies at similar redshifts nor with these high-$z$ quasars passively evolving into the hosts of the low-$z$ quasars.  \\subsection{Host galaxy morphologies compared}  We also compare the morphologies of our quasar hosts with those of the other samples of high-$z$ objects for which there are good estimates of scale sizes in the rest-frame $\\sim V$ or $B$. Though our determinations of scale sizes are very approximate, in no cases did we find evidence for a half-light radius more than 4 kpc ($\\sim$0\\farcs5 at these redshifts). These scale sizes are more consistent with those of the Lyman break galaxies than with those of the high-$z$ radio galaxies or low-$z$ quasar hosts, and our simulations have strengthened this result.  The $z = 2 $ -- 3 HDF-N Lyman break galaxies are found to be quite compact in the rest-frame UV WFPC2 imaging, with scale sizes of $\\sim$0\\farcs2--0\\farcs3 (Giavalisco et al. 1996; Lowenthal et al. 1997). The NICMOS F160W imaging, most comparable to our studies of the MZZ quasars, revealed similar compactness in the LBG sample, with identical or even more compact half-light radii than those measured in the optical (Dickinson 2000). On the other hand, the low-$z$ quasar hosts are found to have half-light radii of 8 -- 15 kpc in the rest-frame $V$ WFPC2 imaging, consistent with their being relaxed giant ellipticals or evolved disk galaxies (McLure et al. 1999; Bahcall et al. 1997).  Comparison to the morphologies of the high-$z$ radio galaxies is less simple, as there may be some dependence of the half-light radii on the luminosity of the radio sources (e.g. Roche, Eales \\& Rawlings 1998; Lacy et al. 2000). Generally, however, at $z \\lesssim 1.5$, the galaxies (with a range in radio luminosities) are found to be well-fit by $r^{1 \\over 4}$ laws, with half-light radii in the range of 5 -- 11 kpc (Roche et al. 1998; Best et al. 1998; McLure \\& Dunlop 2000; Lacy et al. 2000). Pentericci et al. (2000), however, indicate that the powerful radio galaxies in their sample at $z > 2$ have scale sizes 3 times less than those at $z \\sim 1$, while Lacy et al. (2000) find less evolution in scale size. While the best determinations of scale size at rest-frame $V$ are these NICMOS observations of the MRC/USS sample of Pentericci et al. (2000), their sample is fairly small and reveals a large range in scale size. In 5 cases, they were able to fit de Vaucouleurs profiles, and determined $r_{1 \\over 2}$ values ranging from 0\\farcs2 to 1\\farcs6. In Lacy et al. (2000), a summary of the available morphological data in the literature (including the measurements of Pentericci et al.), plus the results of ground-based scale size determinations on the 7C-III sample, indicate that at $z \\sim 2$ -- 3 the mean scale size of high $z$ radio galaxies is $\\sim$0\\farcs6, with the same large range found in the MRC/USS sample. Thus, although some of the high-$z$ radio galaxies have scale sizes as compact as the LBGs or the MZZ quasar hosts, the mean $z \\sim 2 $ --3 radio galaxy is probably still more extended.  With such a small number of objects we cannot effectively judge how significantly ``disturbed'' the morphologies in this sample are in comparison to other samples of high-$z$ galaxies. Three of our 5 quasars have discrete companion objects or a compact component within 10 kpc of the quasar and otherwise seem to have regular morphologies. Overall, this level of ``disturbance'' in the sample is consistent with the complex morphologies seen in the Lyman break and radio galaxy samples. However, even the low-$z$ quasar hosts have some discrete residual components and tidal tails that are visible after subtraction of the main galaxy component (e.g. McLure et al.  1999), and an unusually high incidence of nearby companions is also noted (Bahcall et al. 1997).   \\subsection{Black hole masses, quasar luminosities, and quasar host properties}  We also wish to investigate the relationship between the quasar host properties, the luminosities of the quasar nuclei, and the masses of the nuclear black holes. In particular, a number of results have lent support to the idea that the AGN luminosity and the mass of the host galaxy may be linked. Studies of nearby spheroids indicate that nearly every present-day stellar spheroid contains a massive black hole (BH), and that the masses of the black holes and the luminosities of the bulges (and therefore by inference the masses of the bulges) are linearly correlated (Kormendy \\& Richstone 1996; Magorrian et al. 1998; van der Marel 1999). The stellar velocity dispersions of these bulges may provide a more direct measure of their masses, as evidenced by the even tighter correlation of velocity dispersion with black hole mass (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a). The BH masses determined from this velocity dispersion correlation are consistent with those determined from reverberation mapping of AGN (Gebhardt et al. 2000b). All of these results imply a close link between the formation of the bulges of the galaxies and the formation of the black holes found within them.  If more massive black holes translate into more powerful quasar nuclei, we might then expect that the brighter AGNs (including both brighter radio sources and quasars of greater nuclear luminosity) would have the most luminous and most massive hosts. In fact, there is an upper bound to the luminosity of the low-$z$ quasars that is consistent with Eddington-limited accretion onto supermassive black holes with masses given by the local $M_{BH}$ {\\it vs.} $M_{bulge}$ relation (McLeod, Rieke, \\& Storrie-Lombardi 1999; McLure et al. 1999; Laor 1998).  To make an estimate of the masses of the black holes associated with these MZZ quasars, we can apply the local $M_{BH}$ {\\it vs.} $M_{bulge}$ relation to our host galaxies. We have used the van der Marel (1999) normalization of this relationship, and have made corrections for the much lower M/L ratio at the redshifts of our objects than at $z = 0$ (estimated from the passive evolution models discussed earlier). This results in black hole masses mostly around 1 $\\times$ 10$^{8}$ M$_\\odot$, but ranging from 0.3 -- 2 $\\times$ 10$^{8}$ M$_\\odot$. These values are given in Table \\ref{bhtab}. We then use these derived black hole masses and the $M_B$ nuclear magnitudes of the quasars to estimate the ratio of the quasar luminosity to the Eddington luminosity, making a bolometric correction to the $B$ band of a factor of 12 (Elvis et al. 1994). We find accretion rates that vary from 0.2 to 1.4 times the Eddington limit, with most at 70\\% of the Eddington accretion rate (also in Table \\ref{bhtab}). This is consistent with schemes in which high-$z$ radio-quiet AGN are associated with moderately massive black holes and accrete at close to the Eddington rate, while luminous radio-loud objects are believed to be accreting at sub-Eddington rates and to be associated only with the most massive black holes ($>$10$^{9}$ M$_{\\odot}$) (Rawlings \\& Saunders 1991; McLure et al. 1999; Willot et al. 1999; Lacy, Ridgway \\& Trentham 2000). In contrast, Rix et al. (1999) found comparable host magnitudes for a much brighter set of quasars, and hence found accretion rates that were super-Eddington.  Another way of estimating the black hole masses in AGN is by using the velocity dispersion of the $H\\beta$ emitting clouds in the broad-line region (BLR), (as measured by the $H\\beta$ FWHM), in combination with some estimate of the size of the BLR emitting region, and assuming that the dispersion is from virialized cloud motion (e.g. Laor 1998; Wandel et al. 1999). Wandel et al. compare reverberation mapping estimates of the black hole masses in a sample of Seyferts and quasars to estimates from the $H\\beta$ FHWM, and found they correlated well. Laor (1998) found that the $H\\beta$-derived BH masses in the Bahcall et al. quasar sample correlated well with the black hole masses derived from the Magorrian et al. (1997) $M_{bulge}$ / $M_{BH}$ relation. From the identification spectra for the MZZ quasars (Vitelli et al. 1992), we can estimate very rough CIV FWHMs for 4 of the quasars (in the range 6500 to 7500 km s$^{-1}$), and use these as representative of the $H\\beta$ widths. Laor (1998) estimates the radius of the BLR as proportional to the square root of the bolometric luminosity, and adopting his formulation and our previous estimates of the bolometric nuclear luminosity we find black hole masses for these quasars of about 3 to 8 $\\times$ 10$^{8}$ M$_\\odot$. The derived masses are roughly proportional to those we derived from the van der Marel (1999) relation but are about a factor of 4 larger. However, both Wandel et al. and Laor found that their correlations required separate calibration, and as we have used a different line, and the available spectra are insufficient to make reliable estimates of the line FWHMs, this discrepancy is not surprising.  \\subsection{Implications}  Observations of galaxies at high-redshift to date have focussed largely on two very distinct populations: powerful radio-loud AGN (chiefly radio galaxies, but including radio-loud quasars), and normal star-forming galaxies (the Lyman break population). One of the major goals of our investigation is to better understand the relationship of the hosts of typical radio-quiet AGN to these other populations.  We have found that the radio-quiet quasar hosts are similar to the Lyman break galaxies in terms of rest-frame optical luminosities and sizes, but are considerably less luminous and smaller than radio galaxies at similar redshifts. These results have some interesting implications. We will consider the comparisons to the radio-loud AGN hosts first, and then briefly discuss the Lyman break galaxies.  \\subsubsection{Evolution of quasar and radio galaxy hosts}  Our results reveal an interesting discrepancy between the evolution of the radio-quiet and radio-loud host galaxies. At low-$z$, luminous radio-quiet quasars are found primarily in massive early-type galaxies with luminosities of several times a present-day $L_*$ galaxy (e.g. Bahcall et al. 1997; McLure et al. 1999; McLeod et al. 1999), and have properties that are comparable to those of the radio galaxies. McLure et al. have compared carefully selected samples of radio-quiet quasars, radio-loud quasars, and radio galaxies at $z \\sim 0.2$ and shown their host properties are all similar (with the radio-loud objects in slightly brighter hosts that those of the radio-quiet). However, despite being nearly indistinguishable from radio galaxy and radio loud-quasar hosts at low $z$, at high $z$ the radio-quiet quasar hosts are several magnitudes fainter than the radio galaxies.  The cosmic evolution of the population of powerful radio galaxies has been well-documented over the past decade (e.g.  R\\\"ottgering, Best, \\& Lehnert 1999). Radio galaxies (over a wide range in radio power) are found in relaxed, massive elliptical hosts up to fairly high redshift ($z \\sim 2.5$) and have $K$ magnitudes with low dispersion, consistent with their having formed at much higher $z$ and then subsequently passively evolved into present-day giant ellipticals (Lilly \\& Longair 1984, Lilly 1989, Rigler et al. 1992, Eales et al. 1996, Best et al. 1998, McCarthy 1999). Recent work has indicated that at $z \\sim 2.5$ or higher, radio galaxies have more unsettled morphologies (Pentericci et al. 1999, van Breugel et al. 1999), and there is a larger spread in absolute magnitude (Lacy et al. 2000), indicating that at $z \\sim 3$ -- 5 we may be reaching the epoch of radio galaxy formation. Much less is known about the evolution of the population of the hosts of radio-loud quasars, but the available data paint a broadly similar picture (e.g. Lehnert et al. 1992, 1999; Ridgway \\& Stockton 1997; McLure et al. 1999).  While the most powerful radio galaxies are extremely rare objects, the lower-luminosity radio sources have much higher space densities. For H$_0$ = 50 km s$^{-1}$ Mpc$^{-1}$ and $\\Omega$ = 1, 3CR radio galaxies at $z$ = 2.5 have a co-moving space density of only 0.2 Gpc$^{-3} \\Delta$log$P_{rad}^{-1}$ while the fainter 6C (Eales et al. 1997), 7C (Lacy et al. 1999; Willott 2000), and MRC (McCarthy et al. 1996) radio galaxies have space densities roughly $10^2$ -- $10^3$ times larger (Dunlop \\& Peacock 1990). These values can be compared to the present-day space-density of first-ranked cluster galaxies (roughly 5000 Gpc$^{-3}$, Bahcall \\& Cen 1993). Allowing for a short lifetime for the radio galaxy phase, the evolved descendants of radio sources over a range in luminosities would account for a modest fraction of the first-ranked cluster elliptical galaxies.  In contrast, however, radio-quiet quasars of the luminosity we have studied in this paper are far more common, and therefore are more plausible progenitors of typical present-day early-type galaxies. Radio-quiet quasars with $M_B \\leq$ $-$23 have co-moving space densities at z$\\sim$2 of $\\sim$2$\\times$10$^4$ Gpc$^{-3}$ (Hartwick \\& Schade 1990). Now, for a quasar lifetime on the order of the Eddington growth-time (a few \\% of the Hubble time at z = 2), the implied space density of the present-day descendants is about 10$^6$ Gpc$^{-3}$, which is comparable to the space density of $L_*$ E's and S0's (e.g. Fukugita, Hogan, \\& Peebles 1998). This identification is quite consistent with the correlation between the masses of supermassive black holes and the spheroids in which they live today. A quasar with $M_B$ = $-$24 powered by accretion at the Eddington rate requires $M_{BH} \\sim$ 2 $\\times$ 10$^8$ $M_{\\odot}$, and this supermassive black hole would live today in a spheroid with a mass of about 5 $\\times$ 10$^{10}$ $M_{\\odot}$ and a V-band luminosity of about 1.5 $\\times$ 10$^{10}$ $L_{\\odot}$ $\\sim$ 0.4 $L_*$.  Though inconsistent with passive evolution like that observed in the radio galaxies, our finding L$_*$ hosts at $z \\sim 2$ -- 3 (and the similar results reported by Rix et al. [1999]) are reassuring for standard hierarchical models. The types of quasar hosts we have imaged can easily be the precursors of typical present-day bulge-dominated $L_*$ galaxies, provided that they continue to grow somewhat through mergers from $z \\sim$ 2 to 3 until the present day. Our typical high-$z$ host might need to accrete a factor of few in mass in order to reach the magnitude of a $z = 0$ $L_*$ galaxy.  Indeed, these results agree fairly well with the specific predictions of the hierarchical galaxy formation models of Kauffmann \\& Haehnelt (2000), in which they have addressed both the formation of bulges and the formation and fuelling of their associated black holes. They predict median host luminosities that are somewhat below present-day L$_*$ for quasars at $z=2$ (and even fainter at $z=3$), for quasars with the nuclear magnitudes of our sample. In Figure \\ref{kaufffig} we show our data superimposed on their models. Our galaxies span the range of their model results, though our average host is brighter than the average of the model galaxies shown. In their model, these hosts are still undergoing major mergers which would allow them to evolve into the present day $L_*$ or several $L_*$ galaxies. These models are also able to reproduce the tight correlation of bulge velocity dispersion to black hole mass (Haehnelt \\& Kauffmann 2000).  As the typical quasar lifetime is probably short, in observing active quasars we may be selecting quite different samples of objects at $z \\sim 2$ -- 3, the peak of the quasar number density, from those at low-$z$. At low-$z$, where the number density of quasars has decreased dramatically, the hosts of most quasars are rarer and somewhat more massive than the more common $L_*$ galaxies that are the possible descendants of our high-$z$ quasar hosts. These high $z$ hosts may evolve therefore into more common, less active low-$z$ counterparts, such as Seyferts and quiescent $L_*$ spheroids. This would make their evolution similar to those of the field ellipticals and Lyman break galaxies, (Dickinson et al. 1999; Steidel et al. 1999), whereas the most massive galaxies like the radio galaxies will undergo a very different evolution, associated with an early formation epoch.  \\subsubsection{Relation to the Lyman break galaxy population}  The co-moving space density of the Lyman break galaxies (near the turnover of their UV luminosity function) is roughly 10$^6$ Gpc$^{-3}$ (e.g. Dickinson 1998). If the hosts of radio-quiet quasars are drawn from the Lyman break population, then we conclude that only a few percent of these galaxies need host a radio-quiet quasar at any given time. If there is a population of unbeamed (``type 2'') quasars at high-$z$ (analogous to the Seyfert 2 galaxies at low-$z$) then the total fraction of the Lyman break galaxies hosting an AGN might be several times larger. These relative numbers are consistent with a simple model in which the lifetime of the Lyman break phase (high star-formation rates) is roughly 10$^9$ years (cf. Ferguson, Dickinson, \\& Williams 2000) and the AGN phase is roughly 10$^8$ years (comparable to the characteristic Eddington growth time).  Could a significant fraction (a few percent) of the known Lyman break population population be the ``quasar 2's''? By hypothesis, these would be objects in which the quasar is hidden from direct view along our line-of-sight. By analogy to Seyfert 2 galaxies in the local universe (e.g. Heckman et al. 1997; Gonzalez Delgado, Heckman \\& Leitherer 2000) the rest-frame UV light in these galaxies would be dominated by light from young stars, and so these could naturally enter the color-selected Lyman break samples. If they are like local type 2 Seyferts, they could be distinguished from the majority of ``normal'' Lyman break galaxies because they would show relatively strong but rather narrow ($\\leq$ 10$^3$ km s$^{-1}$) nebular emission-lines. The Ly$\\alpha$ line would normally be the strongest, but lines from highly ionized species like HeII $\\lambda$1640, CIV $\\lambda$1550, and NV $\\lambda$1240 should also be present with strengths of roughly 5 to 20\\% of Ly$\\alpha$. This picture appears to be quite consistent with the spectroscopic properties of the Lyman break galaxies. C. Steidel (private communication) finds that $\\sim$ 1.5\\% of spectroscopically confirmed Lyman break galaxies at z$\\sim$3 are obvious narrow-lined AGNs with Seyfert-2-like spectra. This is a lower bound to the actual fraction, since the weak NV, CIV, and HeII lines could be missed in many other cases.  An obvious way to test the similarity between the Lyman break galaxies and the hosts of high-redshift AGN would be to obtain rest-frame UV images of the quasar hosts in the present sample. Our on-going HST WFPC2 imaging program will make this test possible."
    },
    "0011/astro-ph0011456_arXiv.txt": {
        "abstract": "We measured the K-band luminosity function  using a complete sample of $\\ngal$ morphologically-typed 2MASS galaxies with $\\mu_{K_s}=20$~mag/arcsec$^2$ isophotal magnitudes $7 < K_{20} < 11.25$~mag spread over $\\area$~str.  Early-type ($T \\leq -0.5$) and late-type ($T > -0.5$) galaxies have similarly shaped luminosity functions, $\\alpha_e=-0.92\\pm0.10$ and $\\alpha_l=-0.87\\pm0.09$.  The early-type galaxies are brighter, $M_{*e}=-23.53\\pm0.06$~mag compared to $M_{*l}=-22.98\\pm0.06$~mag, but less numerous, $n_{*e}=(0.45\\pm0.06)\\times 10^{-2}h^3$~Mpc$^{-3}$ compared to $n_{*l}=(1.01\\pm0.13)\\times 10^{-2}h^3$~Mpc$^{-3}$ for $H_0=100 h$~km~s$^{-1}$~Mpc$^{-1}$, such that the late-type galaxies slightly dominate the K-band luminosity density, $j_{late}/j_{early}=1.17\\pm0.12$. Our morphological classifications are internally consistent, consistent with previous classifications and lead to luminosity functions unaffected by the estimated uncertainties in the classifications.  These luminosity functions accurately predict the K-band number counts and redshift distributions for $K \\ltorder 18$~mag, beyond which the results depend on galaxy evolution and merger histories. ",
        "introduction": "The luminosity function (LF) of galaxies, its parameters, dependence on galaxy type, and evolution are fundamental to observational cosmology and the theory of galaxy formation.  Most existing estimates of the luminosity function are based on redshift surveys of galaxies selected from blue photographic plates (CfA/CfA2, Davis \\& Huchra 1982, de Lapparent, Geller \\& Huchra 1989, Geller \\& Huchra 1989, Marzke et al. 1994ab; SSRS2 da Costa et al. 1994, 1998, Marzke et al. 1998; APM, Loveday et al. 1992; ESO Slice, Vettolani et al. 1997, Zucca et al. 1997; 2dFGRS Folkes et al. 1999, Slonim et al. 2000). The luminosity function derivations are usually based on samples of $\\sim 5000$ galaxies.  Blue surveys emphasize galaxies with active star formation, are sensitive to both Galactic and internal extinction, and those based on photographic plates usually have large photometric uncertainties (0.2--0.4~mag). Deep, blue-selected surveys must also include strong, type-dependent K-corrections.  The only ongoing blue survey is the 2dFGRS of 250,000 galaxies.  Most recent surveys have shifted to selecting galaxies in the red, which somewhat reduces the effects of extinction and leads to samples less influenced by recent star formation.  The Century Survey (Geller et al. 1997) used objects selected from red photographic plates with the photometry recalibrated by R$_c$-band drift scans, while the Las Campanas Redshift Survey (hereafter the LCRS, Shectman et al. 1996, Lin et al., 1996, Bromley et al. 1998) selected the galaxies from Gunn~r-band drift scans calibrated to approximate Kron-Cousins R$_c$.  The Sloan Digital Sky Survey (SDSS) will obtain a surface-brightness limited sample to $r'=17.7$~mag with approximately $10^6$ galaxies (see York et al. 2000).  Infrared galaxy surveys have smaller systematic uncertainties than optical galaxy surveys.  They are almost immune to both Galactic and internal extinction, and the K-corrections and luminosity per unit stellar mass are nearly independent of galaxy type (e.g. Cowie et al. 1994, Gavazzi, Pierini \\& Boselli 1996). The determination of infrared luminosity functions has proceeded slowly, however, because of the difficulty of obtaining large complete samples.  Mobasher, Sharples \\& Ellis (1993) and Loveday (2000) obtained infrared photometry of optically-selected galaxies to estimate the infrared LF. Glazebrook et al. (1994, 1995), Gardner et al. (1997) and Szokoly et al. (1998) used relatively deep IR surveys of small regions, where the faintness of the targets makes it difficult to obtain redshifts of the full sample.  De Propris et al. (1998) and Andreon \\& Pello (2000) have also estimated the infrared luminosity function by constructing volume limited samples of galaxies in the Coma cluster. The resulting samples are typically 10 times smaller than published optical samples ($\\sim500$ rather than $\\sim 5000$ galaxies).  The 2MASS project (Skrutskie et al. 1997) is obtaining a complete infrared map of the sky, with a limiting magnitude for its galaxy catalog of $K_s \\simeq 13.5$~mag.  Since 2MASS overlaps the existing optical surveys, it is easy to rapidly generate large, complete infrared redshift surveys.  In this paper we discuss an infrared redshift survey overlapping the CfA2 survey and the updated Zwicky catalog (UZC, Falco et al. 1999).  To a magnitude limit of $K_{20}\\leq 11.25$~mag, $\\sim90$\\% of the galaxies already have redshifts and the remainder were obtained as part of our redshift survey. For the first time we can derive infrared luminosity functions from samples of comparable size to that of the published optical luminosity functions.  In \\S2 we discuss the sample selection, in \\S3 we derive the luminosity function by galaxy type, and in \\S4 we compare the results to other estimates of the luminosity function. In \\S5 we use our luminosity functions to predict the properties of fainter infrared galaxy samples, and in \\S6 we summarize our results. ",
        "conclusions": "We have derived the first local infrared galaxy sample whose statistical uncertainties are comparable to those of local optical galaxy luminosity functions.  We derived both total and morphologically-typed luminosity functions.  Our morphological types are self-consistent (see Kochanek et al. 2000) and our luminosity functions are insensitive to random errors in the classifications and the parameters change as expected when we shift the boundary between early and late-type galaxies.  Like morphologically-typed optical surveys (CfA and SSRS2), we find that the luminosity functions of early and late-type galaxies have similar shapes, $\\alpha\\simeq -0.9\\pm0.1$, in marked contrast to spectrally-typed optical surveys (ESO Slice, LCRS, 2dFGRS) which usually find that the slope steepens for late-type galaxies.  Note, however, that in Kochanek et al. (2000) we find that the spectral classification methods are not self-consistent because of the aperture bias created by using a spectroscopic aperture that is much smaller than the galaxies being observed.  We used galaxies found in both the optical redshift surveys and the 2MASS survey to estimate the magnitude differences between 2MASS and the optical surveys and also between the different optical surveys.  In all surveys, later type galaxies have bluer optical to infrared colors, but the magnitude differences cannot fully explain the discrepancies between the magnitude scales of the luminosity functions.  Our luminosity functions successfully predict the properties of fainter infrared samples until $K \\gtorder 18$~mag where the models have a significant dependence on galaxy evolution and merging histories and the comparison data is probably affected by sample variance.  These results are preliminary, and the sample is still growing rapidly.  In particular, the survey area complete to the current magnitude limit continues to expand rapidly, and it is easy to build complete, deeper samples in restricted areas to extend to fainter absolute magnitude limits.  With complete sky coverage we can use the 2MASS catalog to probe the relative completeness of redshift surveys and to improve our comparisons between the survey magnitude scales.  By combining this with the surface photometry available for 2MASS galaxies we can quantitatively explore the effects of surface brightness selection effects (e.g. Disney 1976, Sprayberry et al. 1997, Dalcanton et al. 1997, Huchra 1999, Cross et al. 2000) on large redshift surveys.  As the coverage gaps are eliminated we can look at density dependences to galaxy properties.  In Pahre et al. (2000) we derive an improved galaxy velocity function ($dn/d\\log v$ instead of $dn/dM$) based on the Tully-Fisher and Faber-Jackson relations derived from the same 2MASS photometry used to derive the luminosity function.  The velocity function is useful because it determines the optical depth of the universe to gravitational lensing (see Falco, Kochanek \\& Munoz 1998) and can be used to probe the evolution of galaxies (see Gonzalez et al. 2000)."
    },
    "0011/astro-ph0011383_arXiv.txt": {
        "abstract": "The observed spectra of the microlensed sources towards the Galactic bulge may be used as a tool for studying the kinematics and extinction effects in the Galactic bulge. In this paper, we first investigate the expected distribution of the microlensed sources as a function of depth within the Galactic bulge. Our analysis takes a magnitude limited microlensing survey into account, and includes the effects of extinction. We show that, in the current magnitude limited surveys, the probability that the source lies at the far side of the bulge is larger than the probability that the source lies at the near side. We then investigate the effects of extinction on the observed spectra of microlensed sources. Kurucz model spectra and the observed extinctions towards the Galactic bulge have been used to demonstrate that the microlensed sources should clearly show the effects of extinction which, in turn, can be used as a statistical measure of the contribution of the disk lenses and bulge lenses at different depths. The spectra of the microlensed sources provide a unique probe to derive the radial velocities of a sample which lies preferentially at the far side of the Galactic bulge. The radial velocities, coupled with the microlensing time scales, can thus be useful in studying the 3-dimensional kinematics of the Galactic bulge. ",
        "introduction": "The topic of gravitational microlensing has experienced a dramatic increase in scientific interest over recent years. This has been largely due to the realization of its wide-ranging applications, such as the detection of planets and the study of Galactic structure. More than 400 microlensing events have now been detected toward the Galactic bulge and Magellanic Clouds by the microlensing survey teams EROS, MACHO, OGLE, DUO, and MOA. The contribution of the bulge stars to the observed microlensing optical depth towards the Galactic bulge was estimated to be approximately 60\\%, the other 40\\% being due to the stars in the Galactic disk \\citep{kir94}. The optical depth to gravitational microlensing, as estimated from the observed microlensing events by OGLE \\citep{uda94}, MACHO \\citep{alc97}, EROS \\citet{afo03}, and MOA \\citet{sum03}, is significantly larger than the values predicted by theoretical models \\citep{pac91}. It has been argued \\citep{pac94} that the observed optical depth can be best explained if the effect of the Galactic bar and its inclination are correctly taken into account. A consequence of the fact that a large fraction of the events are due to bulge-bulge lensing is, as we will show in more detail later, that the lensed stars will preferentially be located on the far side of the bulge in order that there be sufficient stars along the line of sight to cause microlensing. It was suggested by \\citet{sta95} that this would mean that there should be a systematic offset in the apparent magnitude between observed stars and lensed stars. We present a similar discussion by examining the model spectra of different spectral classes and investigating the effects of extinction on stars located on the far side of the bulge.  We first calculate the contributions of the different layers in the bulge to the microlensing optical depth, as has previously been calculated by \\citet{kir94} and \\citet{zha95}. We have, however, considered the additional effect of extinction in some detail. We then proceed with a discussion on how the spectra of the microlensed sources can be used as a measure of the extinction which, in turn, can be used as a statistical measure of the contribution of the disk lenses and bulge lenses at different depths. We show that the spectra of the microlensed sources can be a useful probe for the 3-dimensional structure of the Galactic bulge. We have also undertaken such a spectroscopic study of the microlensed sources, the details of which are the subject of another paper \\citep{kan03}. ",
        "conclusions": "Various Galactic models have been used to demonstrate the theoretical effects of extinction upon the microlensing optical depth. The exponential density model provides the best approximation of the microlensing optical depth when taking extinction into account. For the exponential model, it is shown that stars that are lensed within Baade's window are about 13 times more likely to be at the far side of the bulge than the near side. The estimated optical depth is $\\approx 4.50 \\times 10^{-6}$ for $a_1 = 2.0$ and $a_2 = 3.0$, which is comparable to the observed optical depths of $(3.3 \\pm 1.2) \\times 10^{-6}$ and $(3.9_{-1.2}^{+1.8}) \\times 10^{-6}$ measured by \\citet{uda94} and \\citet{alc97} respectively.  The calculations presented in this paper do not necessarily assume that there is a onoe-to-one relation between distance and reddening. Rather we assume that the sources will be more reddened than the lens stars and, statistically, other non-microlensed sources in the same narrow field. The follow-up study presented in \\citet{kan03}, for example, selects comparison stars which are very close to the microlensed sources so that the angular distance between them, and hence the variation in foreground extinction, is at a minimum.  The effects of extinction on spectra of microlensed sources have been simulated using Kurucz model spectra. It is shown that the majority of microlensed sources should exhibit extinction between $E_{B-V} = 0.4$ and $E_{B-V} = 0.8$. It is also shown that using extinction as a distance indicator for microlensed sources may be used to statistically estimate the distance to the lens from the probability distribution corresponding to the appropriate Galactic model.  The measured extinctions for the microlensed sources can be used to determine the expected stellar number density gradient $a$ which, in turn, can be used to determine the optical depth more accurately. If there is a smooth velocity gradient within the Galactic bulge, one would expect a statistical correlation between the radial velocity and the extinction of microlensed sources which, given enough samples, will provide useful information regarding the 3-dimensional velocity structure of the far side of the Galactic bulge. Given the variation in extinction through the Galactic bulge, the sample would require $\\sim 100$ spectra of microlensed sources in order to clearly show the described effects."
    },
    "0011/astro-ph0011510_arXiv.txt": {
        "abstract": "We present HST NICMOS photometry of the resolved stellar population in the dwarf irregular galaxy NGC~1569. The color-magnitude diagram (CMD) in the F110W and F160W photometric bands contains $\\sim$2400 stars with a formal photometric error $\\lsim\\,0.1$~mag down to \\mj~$\\approx 23.5$ and \\mh~$\\approx 22.5$. The fiducial photometry has a completeness factor higher than 50\\% down to \\mj~$\\approx 21.5$ and \\mh~$\\approx 20.0$. We describe the data processing which is required to calibrate the instrumental peculiarities of NICMOS. Two different packages (DAOPHOT and StarFinder) for PSF-fitting photometry are used to strengthen the photometric results in the crowded stellar field of NGC~1569.  The resulting CMD is discussed in terms of the major evolutionary properties of the resolved stellar populations. For a distance modulus of $(m-M)_0=26.71$ and  a reddening of $E(B-V)=0.56$, our CMD samples stars down to $\\sim$\\,0.8 \\MSUN, corresponding to look-back times of more than 15~Gyr (i.e., an entire Hubble time). This is clear indication of star-formation activity in NGC~1569 spanning an entire Hubble time. The metallicity of the reddest red giant branch (RGB) stars is in better agreement with $Z=0.004$ as measured in HII regions, than with  $Z=0.0004$ as expected from the stellar ages. The presence of --- yet undetected --- very metal-poor stars embedded in the stellar distribution around \\mj~=~22.75 and \\mj~--~\\mh~=~1.15 is, however, not ruled out. The youngest stars ($\\lsim$\\,50~Myr) are preferentially found around the two central super star clusters, whereas the oldest population has a more uniform spatial distribution. A star-formation rate per unit area of 1~M$_{\\odot}$\\,yr$^{-1}$\\,kpc$^{-2}$ and a mass formed in stars of $\\sim1.4 \\times 10^6$~\\MSUN~ in the last 50~Myr are derived from the CMD. The near-infrared (NIR) CMD places strong constraints on the lower limit of the onset of star formation in NGC~1569.  The exceptionally high crowding in the NICMOS images of NGC~1569 is a challenge for the photometric analysis. As a result, optical and NIR images of NGC~1569 sample different populations and cannot be cross-correlated. Nevertheless, we demonstrate the consistency of the star-formation histories derived from the optical and NIR CMDs. ",
        "introduction": "Dwarf irregular and blue compact galaxies are important for our understanding of galaxy formation and evolution. In the models of hierarchical clustering (Kauffmann et al. 1993; Cole et al. 1994), low-mass dwarfs are the first structures to collapse from small-scale primordial density fluctuations in the early universe, giving birth to larger stellar systems through merging. The low metallicity (Kunth 1985; Thuan et al. 1994) and high gas content (Thuan \\& Martin 1981) of nearby dwarf star-forming systems suggests a rather low level of their chemical evolution, and makes some of them (e.g., I~Zw~18 or SBS 0335\\,--\\,052) plausible counterparts to primeval galaxies in the early universe (Izotov \\& Thuan 1999). Starbursting systems could represent the local examples of the faint blue galaxies found in deep imaging surveys as an excess with respect to standard models (e.g., Tyson 1988; Babul \\& Ferguson 1996; Williams et al. 1996). Recent analyses have shown that the excess counts are possibly due to the intrinsically bluest, most compact systems, undergoing strong bursts of star formation (SF) at redshift $z\\gsim0.1$ (e.g., Lilly et al. 1995). Therefore, understanding how the SF proceeds in blue compacts and irregulars (see Tosi 1999 for a review on the different SF regimes) is not only important in its own right but has general implications for astrophysical and cosmological issues.  The presence or absence of an underlying older population in dwarf systems with current SF is crucial for our understanding of galaxy evolution. Did all dwarf galaxies form stars already 10~--~15~Gyr ago and have they experienced discontinuous and/or low-level SF activity since then (Thuan 1983; Legrand 2000)? Have the metal-poorest systems just recently ($\\lsim$\\,100 Myr ago) started forming stars (Searle \\& Sargent 1972; Searle, Sargent, \\& Bagnuolo 1973; Izotov \\& Thuan 1999)? If an old stellar population is present in all dwarf star-forming systems, we would have an indirect validation of the hierarchical collapse model. Identifying local dwarf starbursting galaxies as counterparts to the excess faint blue galaxy population at intermediate redshift would be more straightforward. However, the SF rate (SFR) would have to be at least 1~\\MSUN~yr$^{-1}$ at z$\\approx$1 in order to make the brightness consistent with that of the intermediate redshift population (Babul \\& Ferguson 1996). Primordial systems in our local universe could possibly be identified with HI high-velocity clouds (HVCs). These have been suggested by Blitz et al. (1999) to be just gas contained within dark-matter mini-halos, moving along filaments and accreting onto the Local Group. A physical mechanism would instead be required for dwarf systems to justify the onset of the SF regime only in relatively recent epochs (i.e. Babul \\& Rees 1992). These ``young'' galaxies clearly could not contribute to samples at higher redshift ($z\\gsim0.1$).  The best way to search for old stellar populations is to observe the parent galaxies in the infrared (IR), where the reddening effects are much less severe than in the optical, and to derive the corresponding CMDs, where the old low-mass stars in the RGB evolutionary phase are more easily visible and distinguishable from the younger, more massive objects. The brightest feature of old stellar populations on the CMD is the tip of the RGB (TRGB), which is located at a constant luminosity of $\\log L/L_\\odot\\simeq 3.4$ for ages beyond $\\simeq$\\,1~Gyr (Sweigart, Greggio, \\& Renzini 1990). All the SF prior to this epoch contributes stars at the TRGB, producing a jump in the LF of the red stars. Thus, the detection of the TRGB unambiguously shows that the parent galaxy harbors stellar populations older than 1~Gyr. The non-detection could imply only recent SF (i.e., ages younger than 1~--~2~Gyr), although the TRGB feature may be difficult to detect in the presence of a strong intermediate-age population.  In this paper we report on our NIR observations of the resolved stellar population in the core of the dwarf starbursting galaxy NGC~1569 (= UGC~3056 = Arp~210 = VII~Zw~16 = IRAS~4260\\,+\\,6444). This object was highlighted by Gallagher, Hunter, \\& Tutukov (1984) as an outstanding object in their sample of actively star-forming galaxies. It is a prime candidate for a detailed study of the SF history in a starburst galaxy, since at $D=(2.2\\pm0.6)$~Mpc it is the closest starburst system outside the Local Group (Israel 1988). Some of the properties of NGC~1569 are typical of a dwarf galaxy. With a distance modulus of $(m-M)_0 = 26.71$, its total absolute $B$ magnitude is $M_{\\rm{B,0}}\\approx -17$ (Israel 1988), which is intermediate between the two Magellanic Clouds. Total mass, hydrogen content, and metallicity are estimated to be $M \\approx 3.3 \\times 10^8$~\\MSUN, $M_{\\rm H} \\approx 1.3 \\times 10^8$~\\MSUN~ (Israel 1988), and $Z \\approx 0.25$~Z$_{\\odot}$, respectively (Calzetti, Kinney, \\& Storchi-Bergmann 1994; Gonz\\'alez-Delgado et al. 1997). Thus, NGC~1569 is a gas rich system in a relatively early stage of chemical evolution.  NGC~1569 hosts the two ``super star clusters'' (SSCs) NGC~1569-A (consisting of the two sub-clusters NGC~1569-A1 and NGC~1569-A2) and NGC~1569-B. These high-density stellar clusters were first detected and discussed by Arp \\& Sandage (1985) and Prada, Greve, \\& McKeith (1994) and, more recently, by Gonz\\'alez-Delgado et al. (1997). HST studies were done by O'Connell, Gallagher, \\& Hunter (1994), Ho \\& Filippenko (1996), De~Marchi et al. (1997), Buckalew et al. (2000), Hunter et al. (2000), and Origlia et al. (2000). The SSCs reflect a recent starburst in which at least $10^5$~\\MSUN~ of gas was transformed into stars within about 1~pc. SSCs are similar to young Galactic globular clusters observed shortly after their formation (Melnick, Moles, \\& Terlevich 1985; Meurer 1995; but see Zhang \\& Fall 1999).  The recent strong SF activity around the SSCs of NGC~1569 is also found in multiwavelength studies of the different phases of the  interstellar medium (ISM). The cool gaseous phase is anticorrelated with the recent SF. The HI distribution shows a hole with a diameter of 250~pc, centered on SSC A and encompassing SSC B (Israel \\& van Driel 1990; Stil \\& Israel 1998). The HI, as well as the CO, is preferentially found in the proximity of the most extended HII region of NGC~1569, 90~pc to the west of SSC A (Young, Gallagher, \\& Hunter 1984; Greve et al. 1996; Taylor et al. 1999). Deep H$\\alpha$ images indicate that the warm ionized gas is structured into filaments and arcs with lengths of several kiloparsecs, oriented along the minor axis of the galaxy (Hunter, Hawley, \\& Gallagher 1993; Devost et al. 1997). Typical velocities around 50~--~100~km~s$^{-1}$ and maximum values as high as 200~km~s$^{-1}$ have been attributed to expanding superbubbles with dynamical ages of 10~Myr (Tomita, Ohta, \\& Saito 1994; Heckman et al. 1995). Observations of the hot gas (10$^7$~K) with ROSAT and ASCA (Heckman et al. 1995; Della Ceca et al. 1996)  show that $\\sim$~60\\% of the soft X-ray emission in NGC~1569 is centered in between the two SSCs. Its origin is attributed to thermal emission from the hot gas of a galactic superwind emanating from the disk. The X-ray emission as well as the H$\\alpha$ shells are located in the region where  HI and CO  are absent, another indication of the association of the hot and warm gas  with a galactic outflow.  Observations of the  field stellar population of NGC~1569 are required to trace the SF activity over look-back times greater than 10~Myr. The field stellar population (i.e., outside the unresolved SSCs and HII regions) has been resolved with HST by Vallenari \\& Bomans (1996, VB) on WFPC $V$ and $I$ images, and by Greggio et al. (1998, G98) on WFPC2 $B$ and $V$ images. VB demonstrated that the galaxy experienced a strong starburst from 0.1~Gyr to 4~Myr ago by comparing isochrones to the derived CMD. The well populated asymptotic giant branch (AGB) definitely argues for previous SF in NGC~1569. The authors conclude that a prior SF episode  between 1.5 and 0.15~Gyr ago occurred at a significantly lower rate than the most recent starburst. However, the values of these parameters are uncertain due to observational limitations.  G98 interpreted the data with the synthetic CMD method and showed that in the last 0.15~Gyr the field has experienced a general SF activity of about 0.1~Gyr duration and roughly constant intensity. The SF activity stopped approximately 5~--~10~Myr ago. The estimated SFR during this episode is 4~--~20~\\Myr (depending on the slope of the initial mass function, IMF), or 0.5~--~3~\\Myr~kpc$^{-2}$ for an area of 0.14~kpc$^2$ and a distance of 2.2~Mpc (G98). This value is about 1000 times higher than the current average SFR in our Galaxy and at least two orders of magnitude higher than in Local Group irregulars. A close low-density HI-cloud companion, which is connected to NGC~1569 by an HI bridge, may have triggered the exceptionally high SF activity by interaction (Stil \\& Israel 1998).  The SF history for epochs older than 0.1~Gyr can be studied by searching for AGB and RGB stars in the IR. We thus observed this exceptional starbursting dwarf galaxy with the NICMOS camera on board of HST. The NICMOS field of view was centered on our previous PC field. This paper presents the results of this study. Section~2 describes the data collection and reduction. The photometry applied to the two reduced images in the F110W and F160W filters is in Section~3. Two different PSF-fitting photometry packages were used to assess the reliability of the derived magnitudes. The results of the completeness tests are described. Section~4 contains an analysis of the NIR CMDs and luminosity function (LF). Comparison of these data with stellar evolutionary tracks allows an estimate of the stellar masses, time scales, and metallicities involved in the SF process of this galaxy. Section~5 discusses the cross-correlation between the NIR and optical data of G98. In Section~6 the properties of the resolved stellar population are correlated with the overall starburst activity. Section~7 is a discussion  of all the major factors contributing to making NGC~1569 a challenging case for the photometric reduction. A summary of the major results is in Section~8. ",
        "conclusions": "We have obtained HST NICMOS photometry of NGC~1569 with the goal of studying its intermediate -- old stellar population. Due to the severe crowding conditions and the critical technical features of the NICMOS cameras, we have paid particular attention to the data reduction procedures, using all the most updated methods to overcome the various problems affecting the data. The photometric uncertainties are still high, but the fact of very similar CMDs resulting from two different photometric packages (cf. Figure~\\ref{cmdnic} and Figure~\\ref{figemi}), suggests that the overall statistical properties of the resolved stellar population have been well characterized. The differences in the individual stellar magnitudes of the two alternative photometries (Figure~\\ref{comp}), as well as the artificial star tests (Figure~\\ref{blend}), indicate some difficulties in a good estimate of the background during the PSF-fitting procedure, and actual photometric errors significantly larger than the formal ones.  The severe crowding has prevented us from finding a statistically significant fraction of stars common to both the optical (G98) and NIR (this paper) catalogs, due to a severe blending between blue and red stars spatially coinciding with each other. Nevertheless, we have demonstrated how the optical and NIR CMDs are complementary to each other, and suitable to infer the SFH over the whole look-back time, once the completeness factors are correctly inserted into the simulations.  The NICMOS photometry of NGC~1569 has turned out to constitute an especially challenging case, because of the extreme crowding conditions of this stellar system. We have found the major reason to be the exceptionally high SFR over a long look-back time ($\\tau\\la$\\,0.1\\,Gyr). This implies that all starbursting galaxies with a similar SF history could pose the same photometric challenge if observed with HST in the NIR.  The NIR CMDs presented here clearly show that NGC~1569, in addition to the outstanding young population examined in detail by G98, contains a conspicuous old stellar component, with a well populated AGB (already discovered by VB) and a sufficiently delineated RGB. Hence, the stars in the field of this galaxy cover the whole age range from a few Myr to a Hubble time. The metallicity of the reddest TRGB stars is in better agreement with Padua tracks at $Z=0.004$ than with the set at lower metal content ($Z=0.0004$). We cannot exclude the presence of very metal-poor RGB stars, embedded in the stellar distribution of the more metal-rich objects around \\mj~=~22.75 and \\mj~--~\\mh~=~1.15. The youngest ($\\tau<$\\,50~Myr) stars seem preferentially concentrated around the two SSCs, while the older field stellar population of NGC~1569 has probably a more uniform distribution. The SFR per unit area derived from our NIR CMD for the field stars in the most recent SF period (1~M$_{\\odot}$\\,yr$^{-1}$\\,kpc$^{-2}$) is consistent with the value quoted in G98, once completeness factors are taken into account. The estimated total mass converted into stars ($\\sim 1.4\\times 10^6$~\\MSUN) is about a factor of 5 higher than that found in SSC~A.  The intensity of the past SF activity has presumably been quite lower than the recent one (0.5~--~3 M$_{\\odot}$yr$^{-1}$kpc$^{-2}$), since at the latter rate G98 estimated that the system would run out of gas in only $\\sim$~0.7~Gyr. Also Gallagher et al. (1984) have estimated a SFR that is a factor of $\\sim$\\,4 lower than the present one for look-back times greater than 0.5~Gyr. It is also possible that NGC~1569 has not experienced an almost continuous SF, but a bursting one, with active periods separated by quiescent intervals. To evaluate more precisely at what epochs the SF has actually occurred and with what intensity, it will be necessary to apply to these data the synthetic CMD method described by Tosi et al. (1991) for ground-based data and updated by G98 for HST data. The results of this application will be described in a forthcoming paper (Tosi et al. 2000$b$).  As discussed in Section~4, our data at the faintest magnitudes are too incomplete and uncertain to put stringent constraints on the galaxy distance modulus. Nonetheless, it is apparent that we can exclude distance moduli much different from the adopted value of 26.71~mag. For instance, a modulus either 1~mag shorter or 1~mag longer would make the distribution of the theoretical tracks quite inconsistent with the observational CMD.  NICMOS data of the dwarf irregular NGC~1569 clearly reveal that this galaxy contains an intermediate age -- old stellar population ($\\tau\\,>$\\,1~Gyr). This result is consistent with the conclusions reached by many authors on the ages of the stellar populations in dwarf irregular and blue compact dwarf galaxies. Deep imaging in the NIR (Thuan 1983) and fitting surface brightness profiles in the optical (Papaderos et al. 1996$a$, $b$) demonstrated that unresolved blue compact dwarf galaxies usually have an underlying low-surface brightness component with elliptical isophotes and redder colors, suggesting an older evolved stellar population formed prior to the present starburst. This supports the idea that they are not truly primordial galaxies, but older systems undergoing transient and/or low level SF episodes. Their older stellar component has been resolved into single stars only recently thanks to the capabilities of HST. A classical example is I~Zw~18, observed both in the optical with WFPC2 (Aloisi et al. 1999) and in the near-infrared with NICMOS (\\\"Ostlin 2000): in this case a stellar component as old as at least 3~--~5~Gyr has been found. VII~Zw~403 and Mrk~178 (Schulte-Ladbeck, Crone, \\& Hopp 1998; Schulte-Ladbeck et al. 1999, 2000$b$) are other dwarf galaxies where HST has been able to evidence such an intermediate age -- old stellar population, by revealing their red and old RGB tip stars.  The general scenario emerging from all these studies is that nearby dwarf irregulars are not genuine young stellar systems, local counterparts of the primeval galaxies present in the early universe (Schulte-Ladbeck et al. 2000$a$). If we want to find primordial systems in our local universe we have probably to look elsewhere."
    },
    "0011/astro-ph0011257_arXiv.txt": {
        "abstract": "From the group of galaxies HCG62, we detected an excess hard X-ray emission in energies above $\\sim 4$ keV with \\ASCA. The excess emission is spatially extended up to $\\sim10'$ from the group center, and somewhat enhanced toward north. Its spectrum can be represented by either a power-law of photon index 0.8--2.7, or a Bremsstrahlung of temperature $>6.3$ keV. In the 2--10 keV range, the observed hard X-ray flux, $(1.0\\pm0.3)\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$, implies a luminosity of $(8.0\\pm2.0)\\times10^{41}$ erg s$^{-1}$ for a Hubble constant of 50 km s$^{-1}$ Mpc$^{-1}$. The emission is thus too luminous to be attributed to X-ray binaries in the member galaxies. We discuss possible origin of the hard X-ray emission. ",
        "introduction": "Clusters of galaxies are thought to have released a large amount of dynamical energy in their initial collapse phase. During their subsequent evolution, starburst-driven winds, cluster mergers, radio galaxies, and random galaxy motions may have supplied additional heating energy to the intracluster space. Presumably, these processes have generated energetic particles (e.g. \\citet{kang,takizawa}), as evidenced by diffuse synchrotron radio emission from some clusters.  Such energetic particles are expected to produce nonthermal X-rays as well, by Compton-boosting cosmic microwave background (CMB) photons. Long searches for such effects among galaxy clusters have recently revealed two candidates; the excess soft X-ray emission detected with \\EUVE \\citep{lieu,mittaz,bowyer}, and the spectral hard X-ray tail observed with \\BeppoSAX \\citep{fusco,kaastra}. However, the exact nature of these emission components remain unclear.  Groups of galaxies are the poorest class of galaxy clusters. Their thermal emission is limited to energies below $\\sim 5$ keV, because the temperature of their hot intragroup medium is about 1 keV (e.g. \\citet{mulchaey,fuka96}). Therefore, they allow us to search for nonthermal X-ray emission, even with instruments operating below an energy of $\\sim 10$ keV. Here we report the detection of excess hard X-ray emission from the group of galaxies HCG62 with the \\ASCA GIS \\citep{ohashi,makishima}. We employ 90\\% confidence limits throughout this paper, and use the Hubble constant of 50 km s$^{-1}$ Mpc$^{-1}$. Solar abundances refer to \\citet{anders}. ",
        "conclusions": "From the galaxy group HCG62, we have detected the excess X-ray emission with a very hard spectrum, which extends up to more than 10$'$ from the group center and somewhat enhanced at the north region. Although its surface brightness is only $\\sim$20\\% of that of the GIS background, we have confirmed its reality through careful analysis. Below, we discuss the origin and nature of this phenomenon.  An immediate possibility is collection of binary X-ray sources in the member galaxies of HCG62. However, based on the total optical luminosity of HCG62 ($\\sim1\\times10^{11}L_{\\rm \\odot}$; \\citet{carvalho}) and the optical vs. X-ray luminosity correlation among elliptical galaxies \\citep{matsushita}, this contribution is estimated to be at most $4\\times10^{40}$ erg s$^{-1}$, which is an order of magnitude short of the observed luminosity. A second possibility is assembly of faint active galactic nuclei (AGNs) in HCG62. However, the optical evidence for AGNs in HCG62 is moderate \\citep{carvalho}, and we have already subtracted such candidates based on the \\ROSAT image. Any remaining AGNs are estimated to contribute no more than 30--40\\% of the total hard X-ray emission. Yet another possibility is fluctuation of background faint sources. Utilizing the log$N$--log$S$ relation in the 2--10 keV band \\citep{ueda}, this contribution is estimated to be at most $\\sim2\\times10^{-13}$ erg cm$^{-2}$ s$^{-1}$ over the radius of $3'-15'$, which is again too low to explain the data. From these considerations, we conclude that the excess hard X-ray emission cannot be explained by assembly of discrete hard X-ray sources, whatever their nature be.  Considering the loose constraint on the Fe-K line, the excess emission might be of thermal origin from very hot plasmas. Actually, \\citet{buote} described the \\ASCA spectra of HCG62, integrated over a radius of $0'$--$3'$, by a two-temperature plasma model of $kT=0.7$ and $kT=1.5$ keV. We have independently reconfirmed their results. However, our spectra (Figure 3) that are accumulated over the $3'$--$15'$ range requires a temperature of $> 6.3$ keV; the original two-temperature model found by \\citet{buote} gives a very poor fit ($\\chi^2/\\nu = 1.66$). Thus, thermal emission with an ``ordinary'' temperature cannot explain the data. Insignificant detection of excess hard X-rays over a radius of $0'$--$3'$ might be due to poor photon statistics and complex spectra of thermal components \\citep{fuka98} at the center region, spectral change of hard components, and so on. There might still be plasmas much hotter than the escape temperature, as is actually found in starburst galaxies (e.g. \\citet{ptak}). However, within $15'$ of HCG62, there is no bright IRAS sources with the 60 $\\mu$m flux exceeding 2 Jy (SkyView at High Energy Astrophysics Science Archive Research Center). We therefore conclude that the thermal interpretation of the excess hard X-ray emission is unrealistic.  Given the difficulties with the discrete-source and thermal interpretations of the diffuse hard X-ray emission, we regard nonthermal interpretation as most promising. One popular scenario of non-thermal X-ray production is inverse Compton scattering of the CMB photons by relativistic electrons with Lorentz factor $\\gamma\\sim 10^{3-4}$, as has been invoked to explain the excess hard X-ray emission from rich clusters (e.g. \\citet{fusco}). However, we cannot constrain the intragroup magnetic field in HCG62 due to lack of information on the diffuse radio flux. If we assume a representative magnetic field intensity of 1 $\\mu$G, we would observe synchrotron radio emission with a flux density of $\\sim$0.3 Jy. Since such a strong radio emission is not seen from HCG62 (SkyView; NVSS VLA image), the inverse-Compton interpretation holds for HCG62 only if its magnetic field is much weaker than 1 $\\mu$G. The reality of such a weak magnetic field is an open question, and in an apparent contradiction to the generally accepted inter-galactic field strengths of $\\sim 1~\\mu$G \\citep{kronberg}, even though such a condition is suggested by the \\BeppoSAX and \\RXTE observations of some rich clusters of galaxies \\citep{valinia,fusco}.  An alternative interpretation is nonthermal Bremsstrahlung between the thermal gas and subrelativistic particles, as proposed for the hard X-ray emission from Abell~2199, of which diffuse radio flux is quite weak \\citep{kempner} like in the case of HCG62. Let us assume for simplicity that the nonthermal electrons have typical energies of 10--100 keV, and their spatial density distribution is similar to that of the thermal intragroup gas. Then, the nonthermal to thermal luminosity ratio in the 0.5--10 keV band becomes $\\sqrt{10}\\alpha$, where $\\alpha$ is the density ratio of the nonthermal electrons to the thermal ones. The observed luminosity ratio of $\\sim 0.2$ implies $\\alpha\\sim$0.06, indicating that the energy density of nonthermal electrons is 0.6--6 times as high as that of thermal electrons (depending on the spectrum). If this is the case, the mechanism of such particle acceleration becomes an important issue. In addition, the non-thermal pressure associated with such a particle population would considerably increase the total mass of HCG62, and hence its dark-matter content, estimated from the X-ray data.  The diffuse hard X-ray emission has been observed with \\ASCA from some other galaxy groups as well \\citep{fuka99}. However, its prominence relative to the thermal X-ray emission appears to scatter from object to object like rich clusters \\citep{molendi}, with HCG62 one of the strongest case. Although it is yet to be studied what makes such variety, the hard X-ray emission might be related with transient phenomena such as mergers.  The authors are grateful to Dr. Takizawa for helpful discussions, and to an anonymous referee for helpful comments. This research was supported in part by the Grants-in-Aid for the Center-of-Excellence (COE) Research of the Ministry of Education, Science, and Culture in Japan (07CE2002)."
    },
    "0011/astro-ph0011061_arXiv.txt": {
        "abstract": "We present a study of selected potential $\\gamma\\,$Doradus stars observed in the Geneva photometric system. Eleven stars were monitored at least during the last year of functioning of the 70 cm Swiss telescope at La Silla, leading to about one thousand photometric measurements and permitting a detailed analysis of these stars. After the analysis, five stars are thought to be constant stars. We report three new $\\gamma$ Dor stars. A spectroscopic campaign has been started at La Silla with the new 1m20 Swiss telescope, equipped with the spectrograph CORALIE, of which we show the results for the selected stars. ",
        "introduction": "The $\\gamma\\,$Dor stars constitute a new class of pulsating variables that was discovered only recently. Krisciunas (\\cite{Krisciunas}) and Zerbi (\\cite{Zerbi}) review the history of the discovery and the current observational status of this group of variables. We refer to these papers for a general introduction on $\\gamma\\,$Dor stars, but we mention here that these objects are thought to exhibit multiperiodic high-order g-mode pulsations with periods in the range of 0.5 - 3 days.  Much effort is currently made to find new members of this group, in order to constrain their pulsational characteristics and their position in the HR\\,diagram. This is especially relevant since the subject of the precise excitation mechanism is not yet settled. Promising theoretical models are currently being worked out. No direct confrontation between the new theoretical ideas about the driving and the observational status has been performed yet, one of the reason being that the number of {\\it bona fide} members is still low and that these stars have quite a large variety in their observational behaviour.  We already performed a systematic detection and classification of $\\gamma\\,$Dor stars from the Hipparcos Periodic Annex using the Geneva photometric system. The method used was a multivariate discriminant analysis and it led to the discovery of 14 new $\\gamma\\,$Dor candidates (Aerts et al. \\cite{Aerts&al}). Moreover, Handler (\\cite{Handler}), searching for new $\\gamma\\,$Dor stars in the Hipparcos catalogues of the periodic and unsolved variables, found respectively 13 and  18 new ``prime'' $\\gamma\\,$Dor candidates which are non-redundant with other studies.  The paper is organised as follows. In Sect.\\,2, we present the obtained data from the Geneva photometry, and compare it to the Hipparcos photometry. In Sect.\\,3, we search for periodic behaviour in the data. In Sect.\\,4, we present the spectra of the stars. Finally, we end with some conclusions in Sect.\\,5. ",
        "conclusions": "The more thoroughly the candidates were studied the more they were rejected. We end therefore with a shorter list of $\\gamma\\,$Dor stars. Establishing such a list needs strong efforts, but it is important for describing precisely the physical properties of this pulsating star group.  We have found three new $\\gamma\\,$Doradus stars (HD\\,12901, HD\\,48501, HD\\,26298). One of them (HD\\,12901) was indicated by Handler (\\cite{Handler}) to have a possible $\\delta\\,$Scuti behaviour on the basis of Hipparcos data. Our Geneva photometry, however, clearly points towards multiple periods longer than those of $\\delta\\,$Scuti stars and in the region expected for $\\gamma\\,$Doradus stars. HD\\,48501 is also a multiperiodic variable. The line-profile variability of these three new $\\gamma\\,$Dor stars will be discussed in a separate paper.  It is presently important to find as many confirmed $\\gamma\\,$Doradus stars as possible, in order to constrain the suggested theoretical excitation models. In particular, the question remains whether or not the $\\gamma\\,$Doradus phenomenon only occurs for F\\,0 -- F\\,2 type stars or also for much cooler objects as suggested by Aerts et al.\\ (\\cite{Aerts&al}). If the latter is true, then it is tempting to assume that the $\\gamma\\,$Doradus stars may exhibit simultaneously opacity-driven g-modes and stochastically excited solar-like p-modes. The occurrence of both such types of pulsations would broaden the interest for the stars because they would become unique objects for seismology. In this respect, the search for solar-like oscillations in bright selected $\\gamma\\,$Doradus stars by future seismology space missions, such as MOST, MONS and COROT, would be extremely valuable."
    },
    "0011/astro-ph0011127_arXiv.txt": {
        "abstract": "We review recent relativistic hydrodynamic simulations of jets, and their interpretation in terms of the results from linear stability analysis.  These studies show that, interpreted naively, the distribution of synchrotron intensity will in general be a poor guide to the physical state (density and pressure) of the underlying flow, and that even if the physical state can be inferred, it, in turn, may prove to be a poor guide to the source dynamics, in terms of the transport of energy and momentum from the central engine. However, we demonstrate that an interplay of simulation and linear stability analysis provides a powerful tool for elucidating the nature and character of structures that jets may sustain. From such studies we can explain the complex behavior of observed jets, which manifest both stationary and propagating structures, without recourse to ad hoc macroscopic disturbances. This provides a framework for the interpretation of multi-epoch total intensity data wherein an understanding of the character of individual flow features will allow the effects of physical state and dynamics to be deconvolved. ",
        "introduction": "While computational fluid dynamic simulations of astrophysical jets have been performed for two decades, it is only since the mid-1990s that this approach has been applied to non-steady albeit non-MHD {\\it relativistic} flows (van Putten 1993; Duncan \\& Hughes 1994; Mart\\'\\i, M\\\"uller, \\& Ib\\'a\\~nez 1994, Mart\\'\\i\\ et al., 1995, 1997; Komissarov \\& Falle 1998; Rosen et al. 1999). Such studies have a dual significance. Firstly, that the relativistic Euler equations are inherently more nonlinear than their classical cousins, that the Lorentz factor couples the flow along each coordinate direction, and that there is no intrinsic limit to either shock thinness or compression, imply that the distributions of density, pressure and velocity will be quite different from those computed classically. Secondly, such computations provide a physically meaningful picture of the flow dynamics, from which Doppler shift, boost, aberration and time delays may be inferred, allowing the production of simulated source maps that incorporate all the pertinent relativistic effects.  With the parallel advances in the area of VLBI/P mapping, which have led to data sets with hitherto unseen spatial, temporal and spectral resolution, and dynamic range, it might be supposed that, by confronting these new simulations with the current and emerging data, the state and dynamics of extragalactic jets on a range of length scales could now be probed in detail.  The goal of this review is to urge caution on those involved in this enterprise, a warning to the curious (James 1984) on the periphery of the subject, and encouragement that progress can be made, but that such progress is probably contingent on our acquiring a deeper understanding of jet dynamics -- through, for example, carefully integrating analytic studies of jet stability with studies based on CFD simulations.  A good case can be made that, at least on observable scales, the magnetic field does not play a dynamically important role; strong evidence exists for a turbulent field in the quiescent state of both BL~Lacs and QSOs (Jones et al. 1985; Hughes, Aller \\& Aller 1989a,b).  The magnetic field responds readily to the underlying flow, and VLB polarimetry will surely provide an invaluable tool for probing jet curvature, shear and the structure of internal, in general oblique, shocks, through the imprint of these features on the field, and thus on the polarized emission.  However, given the field's likely insignificance as regards the flow dynamics, and the fact that a rigorous calculation of the flow emissivity would require a knowledge of the evolution of the distribution function of the radiating particles as much as a knowledge of the magnetic field (Tregillis et al. 1999; see also Tregillis, Jones \\& Ryu these proceedings), for the present discussion we shall consider purely hydrodynamic flows, and the total intensity only. ",
        "conclusions": "Significantly perturbed jets continue to carry energy and momentum far from their origin, but a) the intensity maps may be a poor guide to the energy and density distributions; b) the energy and density distributions may be a poor guide to the source dynamics. We suggest that the interpretation of total intensity data will be facilitated if it can be done in the context of understanding the internal structures that jets may and do support. Specifically, linear stability analysis aids in the interpretation of CFD simulations, while the latter validate extrapolating stability analysis into the nonlinear regime to predict jet structures. We thus have a powerful set of tools with which to confront the most easily addressed data sets -- those involving slow moving, optically thin features (e.g., as seen in M~87). The CFD simulations suggest that jets subject to particular disturbances may exhibit a few, identifiable modes, and if we can test our theoretical tools on data such as that available for M~87, we will be in a position to predict the characteristics of the modes most likely to influence jet structure, for the more challenging case of rapidly evolving, partially absorbed flows."
    },
    "0011/astro-ph0011194_arXiv.txt": {
        "abstract": "We analyzed RXTE data of GRS 1915+105 of X-ray low states between 1996 and 1998. The X-ray spectrum is dominated by the power law component. We found (i) that the power law is pivoting at 10-30 keV, (ii) that the power law slope correlates with radio flux, (iii) three different groups  of 3-10~keV residuals, (iv) that one of these residual groups correlates with the power law slope. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011531_arXiv.txt": {
        "abstract": "In extended quintessence models, a scalar field which couples to the curvature scalar $R$ provides most of the energy density of the universe.  We point out that such models can also lead naturally to a decrease in the primordial abundance of helium-4, relieving the tension which currently exists between the primordial helium-4 abundance inferred from observations and the amount predicted by Standard Big Bang Nucleosynthesis (SBBN) corresponding to the observed deuterium abundance.  Using negative power-law potentials for the quintessence field, we determine the range of model parameters which can lead to an interesting reduction in the helium-4 abundance, and we show that it overlaps with the region allowed by other constraints on extended quintessence models. ",
        "introduction": "A great deal of observational evidence currently points toward a cosmological model with a nonzero cosmological constant $\\Lambda$.   A combination of the supernova Ia measurements \\cite{SNIa}, measurements of the baryon fraction in galaxy clusters \\cite{WNEF}, and the location of the first acoustic peak of the microwave background anisotropies \\cite{peak} suggests a model with $\\Omega_{{\\rm M}} = 0.3 - 0.4$ and $\\Omega_\\Lambda = 0.6 - 0.7$.  A genuine cosmological constant is a disaster from the standpoint of particle physics, in which the most natural value of $\\Lambda$ is many orders of magnitude larger than the ``observed'' value (see Ref.~\\cite{Wein} for a review). However, even if a plausible particle physics mechanism were developed to produce such a small value of $\\Lambda$, a second problem remains: why are $\\Omega_\\Lambda$ and $\\Omega_{{\\rm M}}$ comparable today? Since $\\Omega_\\Lambda$ and $\\Omega_{{\\rm M}}$ scale very differently with the cosmological expansion factor, this coincidence suggests that we live in a very special epoch.  This problem has been dubbed the ``cosmological constant coincidence problem'' to distinguish it from the more fundamental cosmological constant problem \\cite{ZWS,SWZ}.  A possible solution to the coincidence problem is to assume that the apparent cosmological constant is not, in fact, a true constant vacuum energy density, but instead is due to the energy density of a scalar field $\\phi$, a possibility which has come to be known as {\\em quintessence}. It might then be possible to couple the behavior of this field to the background matter density in such way as to achieve `naturally' the desired result, namely $\\Omega_\\phi \\approx \\Omega_{{\\rm M}}$ today.  Of course, in this case, the scalar field may have an equation of state intermediate between matter and a cosmological constant. This possibility was explored in some detail by Zlatev et al. \\cite{ZWS} (see also reference \\cite{SWZ}) who argued that a certain class of solutions (``tracking solutions'') will evolve toward the desired behavior independent of the initial conditions.  Although such models still require fine tuning to produce $\\Omega_\\phi \\approx \\Omega_{{\\rm M}}$ today \\cite{LS}, they seem to be pointing in a more plausible direction toward this result.  These investigations assumed minimally-coupled scalar fields, but it was soon realized that coupling the scalar field to the curvature scalar $R$ opens up another range of possibilities, dubbed ``extended quintessence.''  Uzan \\cite{Uzan} examined the general case of a non-minimally coupled scalar field evolving in an exponential or power-law potential, and Amendola \\cite{Amendola} explored a general class of couplings and potentials.  Non-minimally coupled models suffer from the potential problem that the gravitational constant $G$ varies in time \\cite{Amendola,Chiba}, but they have nonetheless received a great deal of recent attention as possible models for quintessence \\cite{pbm} - \\cite{bmp}.  In this paper, we point out an interesting consequence of non-minimally coupled quintessence models:  under some circumstances these models can lead to a reduction in the primordial helium abundance.  This is of interest because of a ``tension'' between the recent estimate of the primordial deuterium abundance, $D/H = 3.0 \\pm 0.4 \\times 10^{-5}$ \\cite{omeara}, which leads to an estimate for the baryon/photon ratio of $\\eta = 5.6 \\pm 0.5 \\times 10^{-10}$, and the corresponding SBBN predicted primordial helium-4 abundance of $Y_P = 0.248 \\pm 0.001$. In contrast, the actual primordial helium abundance is likely lower. For example, Olive and Steigman found \\cite{Oliver-Steigman} \\begin{equation} Y_P = 0.234 \\pm 0.003 ({\\rm stat.}), \\end{equation} while Izotov and Thuan obtained a higher value \\cite{Izotov-Thuan} \\begin{equation} Y_P = 0.244 \\pm 0.002 ({\\rm stat.}). \\end{equation}  While this apparent discrepancy is insufficient to discard SBBN, it certainly represents a ``tension'' in the model.  Furthermore, in addition to the three standard model neutrinos, a sterile neutrino may be needed to explain the results from neutrino oscillation experiments \\cite{neutrino}. If either or both of these were confirmed, or if there are any other light particles in the Universe, the breach between theory and observation on $~^{4}\\mathrm{He}$ could become even wider.  In light of these SBBN results, any ``natural'' mechanism which lowers the primordial helium-4 abundance must be regarded as interesting.  In this paper we show that extended quintessence models provide just such a mechanism.  In Ref. \\cite{pbm}, Perrotta et al. investigated the BBN constraint on the extended quintessence model.  They showed that the model is not ruled out by increasing helium, but they did not consider the more interesting possibility of decreasing helium.  In the next section, we review the evolution of the scalar field in these models and calculate the reduction in the abundance of helium-4 as a function of the model parameters.  We determine whether an interesting reduction is consistent with other constraints on extended quintessence. Our results are summarized in Sec. 3.  We shall use units in which $\\hbar=c=8\\pi G_{N}=1$ throughout, unless otherwise stated. ",
        "conclusions": "We have investigated the effect of NMC quintessence model on BBN. In models with $\\xi>0$, the expansion rate is smaller during BBN, and less helium is produced. The amount of this reduction depends primarily on the coupling constant and the slope of the potential, and also very weakly on the initial value of the field.  However, the coupling constant $\\xi$ is limited by solar system experiments. We find that for negative power-law potentials of the form $V = V_0 Q^{-\\alpha}$, the magnitude of the reduction increases with $\\alpha$. This is not a desirable state of affairs from the point of view of reducing primordial helium.  The reason is that negative power-law potentials give an equation of state $w = p/\\rho$ in which $w$ approaches 0 as $\\alpha$ increases.  In Figure 10 is shown the relation between $w$ and $\\alpha$.  For instance, $\\alpha = 6$ corresponds to $w = -0.43$, and a variety of observations argue against a value of $w$ much larger than this (see, e.g., Ref. \\cite{wcos}).  If we take $\\alpha \\le 6$ as a reasonable upper bound on such models, then we find $\\Delta Y \\le 0.0061$.  This is an interesting reduction in primordial helium (from $Y_P = 0.248$ to $Y_P = 0.242$), albeit it requires us to push all of the parameters in the model to the extreme acceptable limits. For example, while a reduction as small as $\\Delta Y \\approx 0.002 - 0.004$ would be sufficient to reconcile the O'Meara et al. deuterium abundance \\cite{omeara} with the Izotov and Thuan helium value \\cite{Izotov-Thuan}, a much larger reduction would be required if the Olive and Steigman \\cite{Oliver-Steigman} helium abundance were adopted.  If there is indeed a breach between the observed helium abundance and the predictions of SBBN it could be healed by this NMC model. Alternatively, a reduction in $Y_P$ of 0.0061 corresponds to a shift in the equivalent number of neutrinos by $\\Delta N_{\\nu} = - 0.47$.   \\begin{figure}[tbp] \\begin{center} \\epsfig{file=alphaw.ps,width=3.5in} \\caption{$w$ versus $\\alpha$, with strongest allowed coupling $\\xi$.} \\label{alphaw} \\end{center} \\end{figure}    We have shown that certain classes of NMC quintessence models lead to a ``natural\" reduction in the primordial helium abundance. Although we have concentrated on the negative power law potentials, our results generalize easily to other potentials.  In particular, one of the major problems with exponential potentials is that they lead to an overproduction of helium \\cite{expQmodel}.  This problem could be ameliorated in the way we have suggested. Of course, the tracker solution for such models generically predicts $w = 0$ at present \\cite{expQmodel} and thus these models have problems matching other observations.  We also expect that these results would be changed if we used a different coupling between $Q$ and $R$ than that in Eq. (\\ref{FQ}). Thus, it might be possible to obtain even larger values for $\\Delta Y$ in such models."
    },
    "0011/astro-ph0011525_arXiv.txt": {
        "abstract": "Several of the models for origin of the highest energy cosmic rays also predict significant neutrino fluxes.  A common factor of the models is that they must provide sufficient power to supply the observed energy in the extragalactic component of the cosmic radiation.  The assumption that a comparable amount of energy goes into high-energy neutrinos allows a model-independent estimate of the neutrino signal that may be expected. ",
        "introduction": "An important argument in favor of supernova explosions as the power source for galactic cosmic rays is the fact that kinetic energy of the ejecta supplies the right amount of power.  The energy content of the cosmic radiation is \\begin{equation} \\rho_E = {4\\pi\\over c}\\,\\int\\,E\\,\\phi(E)\\,dE\\;\\approx\\;10^{-12}\\; {\\rm erg/cm}^3, \\label{Edensity} \\end{equation} where $\\phi(E)=dN/dE$ is the measured local flux of cosmic rays corrected for the effect of solar modulation. In the source region the average energy density in cosmic rays is related to the average production rate per unit volume, $q(E)$, by \\begin{equation} \\rho_E\\;=\\;q(E)\\times\\tau_{esc}(E), \\label{CRsource} \\end{equation} where $\\tau_{esc}$ is the characteristic residence time of cosmic rays in the source region (for example the disk of the galaxy).  The characteristic time, $\\tau_{esc}(E)$, decreases with energy \\cite{Garcia-Munoz} so that the observed spectrum is somewhat steeper than the source spectrum.  Given an estimate of $\\tau_{esc}$ and the rate of supernova explosions, it is possible to estimate the fraction of energy of supernova explosions needed to maintain the galactic cosmic rays in steady state, assuming that supernovae provide the power, $q(E)$.  Assuming a rate of three supernovae per century with a kinetic energy of $10^{51}$ ergs per supernova, the conclusion is that an efficiency of $\\sim 10$\\% for conversion of kinetic energy of supernova ejecta into relativistic cosmic rays would suffice. In the years since this coincidence was pointed out \\cite{Ginzburg}, a theory of cosmic-ray modified shocks with a high efficiency for particle acceleration has been developed \\cite{crmodified}. A compelling feature of this theory is that it produces a spectral index that fits in well with what is observed after energy dependence of propagation is accounted for.  Here we want to apply a similar analysis of energetics to cosmic rays of extragalactic origin  The reason to focus on the high energy end of the cosmic ray spectrum in connection with high energy neutrino astronomy is that several models have been suggested as sources of ultra-high energy cosmic rays which would also be likely sources of high energy neutrinos.  These include some models of active galactic nuclei (AGN) and some models of gamma-ray burst (GRB) sources.  In the case of galactic cosmic rays, if the production rate of neutrinos is proportional to that of cosmic rays, $q_\\nu(E)\\,=\\,f\\times q(E)$, we would expect a flux of neutrinos related to the cosmic-ray flux by \\begin{equation} \\phi_\\nu \\propto {V_D\\over c}\\times f\\times \\phi_{cr}, \\label{CRneutrino} \\end{equation} where $f$ is the efficiency with which the cosmic rays interact to produce neutrinos (either in the source region or in the interstellar medium) and $V_D$ is the velocity of diffusion of cosmic rays out of the galaxy. $V_D\\sim R_{galaxy}/\\tau_{esc}$ depends on the model of cosmic-ray propagation in the galaxy, but in any case $V_D\\ll c$.  Thus the galactic neutrino flux should be suppressed by a large factor relative to the parent cosmic rays because of their straight-line propagation out of the galaxy. For the same reason, neutrinos, like gamma-rays, would be expected to show the structure of the galactic disk.  For a cosmological distribution of sources of ultra-high energy cosmic rays we assume initially, for the rough estimates discussed below, that \\begin{equation} V_D\\;\\sim\\;{R_H/\\tau_H}\\;\\sim\\;{3000\\,{\\rm Mpc}\\over 10^{10} yrs}\\;\\approx\\;c. \\label{Hubble} \\end{equation} This approximation neglects evolution and assumes that the intergalactic magnetic fields are weak enough so that charged particles can reach us from distant sources in less than the age of the universe.  The situation for super-GZK particles, where the maximum distance of propagation is limited to $R\\ll R_H$ by photo-pion production, needs a separate discussion.  The possibility of large inter-galactic magnetic fields and local concentrations of sources \\cite{local1,local2} would also have to be considered for a full treatment of the cosmic-ray spectrum.  We use the approximation of Eq.~\\ref{Hubble} primarily to derive an estimate of the diffuse flux of high energy neutrinos that may be associated with the sources of extragalactic cosmic rays.  The outline of the paper is as follows.  First we discuss the transition from galactic to extragalactic cosmic radiation in order to define an extragalactic component.  Then we compare the power needed for the extragalactic cosmic radiation with that available from various potential sources.  We conclude with a review of the predictions for detection of high energy neutrinos. ",
        "conclusions": ""
    },
    "0011/astro-ph0011239_arXiv.txt": {
        "abstract": "We present our ground-based optical imaging, spectral analysis, and high resolution radio mapping of the gravitational lens candidate FBQ 1633+3134.  This $z=1.52$, $B=17.7$ quasar appears double on CCD images with an image separation of $0\\farcs66$ and a flux ratio of $\\sim3:1$ across $BVRI$ filters.  A single $0.27$ mJy radio source is detected at 8.46 GHz, coincident to within an arcsecond of both optical components, but no companion at radio wavelengths is detected down to a flux level of $0.1$ mJy ($3\\sigma$).  Spectral observations reveal a rich metal-line absorption system consisting of a strong \\ion{Mg}{2} doublet and associated \\ion{Fe}{1} and \\ion{Fe}{2} absorption features, all at an intervening redshift of $z=0.684$, suggestive of a lensing galaxy. Point spread function subtraction however shows no obvious signs of a third object between the two quasar images, and places a detection limit of $I \\gtrsim 23.0$ if such an object exists.  Although the possibility that FBQ 1633+3134 is a binary quasar cannot be ruled out, the evidence is consistent with it being a single quasar lensed by a faint, metal-rich galaxy. ",
        "introduction": "The number of quasars multiply imaged by gravitational lenses has steadily increased in the last decade.  There are now over 40 well-established cases of galaxy-size gravitational lenses (Kochanek et al. 1998).  The usefulness of gravitational lenses are well known (Kochanek \\& Hewitt 1996).  For example, measurements of a time-delay between the light curves of individual images can lead to a determination of the Hubble constant (Refsdal 1964), and the observed frequency of strong lensing events can provide statistical constraints on the cosmological constant (Kochanek 1996).  Also, since the observed image configuration is sensitive to the gravitational potential of the lens, gravitational lensing offers a means to probe the matter distribution of the lensing object (e.g., Maller et al. 2000).  This provides a method to sample the dark matter halos of intermediate redshift galaxies which is independent of the galaxy's luminosity.  The identification of a close separation (i.e., $<10\\arcsec$) double quasar, however, is not confirmation of a gravitational lens.  An alternative explanation is a physical binary quasar (e.g., Kochanek et al. 1999, Mortlock et al. 1999). In principle, the distinction between these two possibilities is clear: either there is direct evidence for an intervening galaxy between the quasar images, confirming the lensing scenario, or the spectral properties of the pair differ to the extent that the most likely explanation is that of two separate quasars.  In the latter case, it is usually the failure to detect one of the quasars at radio wavelengths that eliminates the lensing hypothesis (e.g., Djorgovski et al. 1987; Mu\\~noz et al. 1998). In practice, there are a number of quasar pairs that do not fall neatly into either category (Kochanek et al. 1999).  The astrometric and spectral characteristics of these pairs are typical of gravitational lenses, yet no direct evidence for a lensing galaxy is found.  The double quasar FIRST J163349.0+313412 (hereafter, FBQ 1633+3134), which we analyze here, is best classified as one of these ``border-line'' systems.  In this paper, we present our ground-based optical imaging, spectral observations, and radio mapping of FBQ 1633+3134, and discuss the likelihood that the system is a gravitational lens.  FBQ 1633+3134 (16$^{\\scriptsize{\\mbox{h}}}$ 33$^{\\scriptsize\\mbox{m}}$ 48$\\fs$97, +31$\\arcdeg$ 34$\\arcmin$ 11$\\farcs$8; J2000.0) was originally identified as a $z=1.52$, $B=17.3$ quasar from the FIRST Bright Quasar Survey (FBQS; Gregg et al. 1996).  The target was first detected as a 2 mJy radio source at 20 cm as part of the FIRST radio survey (Becker, White, \\& Helfand 1995), and follow-up spectral observations taken with the Lick 3 m telescope confirmed the radio source as a high-redshift quasar (White et al. 2000).  The FBQS offers a major advantage over fainter quasar surveys for the purpose of discovering lensed quasars.  Since the majority of the FBQS targets are bright ($B<18$), the probability of finding a quasar multiply imaged by gravitational lensing is enhanced over a sampling of fainter quasars due to the magnification bias effect (Turner, Ostriker, \\& Gott 1984). Using the 2.4 m Hiltner Telescope at the MDM Observatory, we have exploited this bias by obtaining optical images of the brightest ($B<18$) and more distant ($z>1$) of the FBQS quasars to search for arcsecond scale multiple imaging indicative of strong gravitational lensing.  One such observing run in 1996 December discovered the quasar FBQ 0951+2635 as a sub-arcsecond separation gravitational lens (Schechter et al. 1998).  A subsequent observing run was carried out during 1997 June and targeted an additional $100$ FBQS quasars in a further search for lenses. During this run, a $3$ minute $R$-band exposure of FBQ 1633+3134 showed two stellar-like point sources separated by $\\sim 0\\farcs7$, and immediate follow-up exposures showed the same separation and similar flux ratios in $BVI$ bandpasses as well, providing the initial evidence that FBQ 1633+3134 might be gravitationally lensed.  We describe our initial optical observations and analysis of FBQ 1633+3134 in \\S2, as well as higher resolution follow-up optical imaging taken immediately afterwards using the 3.5 m WIYN telescope.  In \\S3, we describe the original FBQS spectrum of FBQ 1633+3134, as well as a follow-up spectrum taken with the Keck 10 m telescope.  While radio observations were originally obtained as part of the FIRST survey, we have since obtained higher resolution VLA radio imaging of FBQ 1633+3134 following the optical detection of the object's double nature.  We present these radio observations and their implications in \\S4.  Finally, we discuss lens modeling and interpretation of the system in \\S 5, and \\S 6 summarizes our findings and conclusions for FBQ 1633+3134. ",
        "conclusions": "From an observational standpoint, the requirements for a double quasar to be confidently classified as a gravitational lens can be stated as follows:  First, since gravitational lensing is achromatic, the optical flux ratio of the components should (ideally) be self-consistent across broadband filters.  In practice, the ratios will not agree exactly due to differential reddening through the lens galaxy (McLeod et al. 1998) and microlensing by stars in the lens galaxy (Wisotzki et al. 1993).  More importantly, the radio flux ratio, if available, must also agree with the optical ratio, again because of the achromatic nature of lensing. Second, spectral observations of the quasar pair must support the conclusion that both objects are identical redshift quasars with similar (although not necessarily identical; Wisotzki et al. 1993) emission and continuum properties.  Third, the lensing galaxy itself must be detected, ideally by direct imaging,  although weaker arguments based on the presence of intervening absorption features in the QSO spectrum have been used in the past (e.g., Hewett et al. 1994).  In practice, the first two requirements are necessary for a lens classification, while the third (direct imaging of the lensing galaxy) is usually sufficient.  FBQ 1633+3134 meets the necessary requirements for a gravitational lens. The broadband flux ratio ($\\sim 2.5:1$ in $B$, $\\sim 3.3:1$ in $VR$, \\& $I$) is self-consistent across broadband wavelengths.  Although only one component was detected in radio, the lower limit to the radio flux ratio of $\\gtrsim 3:1$ is not in severe conflict with the corresponding optical flux ratio.  Also, the spectral analysis argues against either component of the pair being a Galactic star, and the unresolved emission profiles from the pair are consistent with two identical-redshift quasars.  The sufficient condition that could elevate FBQ 1633+3134 to the status of a lensed quasar, direct detection of the lensing galaxy, is not realized with our optical imaging.  Subtraction of two PSFs for the WIYN data presented in \\S2 show no significant indication of unaccounted flux in the system, and place a magnitude limit of $I>23.0$ using a simple SIS model for the lensing potential.  Although the hypothesized lens galaxy is not detected directly, a weaker argument for its presence can be made from the spectral identification of a rich $z_{abs}=0.684$ metal-line absorption system, which is within $1\\sigma$ of the median lensing redshift.  Also, the estimated luminosity of the suspected galaxy does not differ strongly from an $L_*$ galaxy for either of the galaxy types considered here.  On the other hand, the predicted apparent magnitude of a late-type spiral at the absorber redshift is $\\sim2$ magnitudes above the $I$-band detection threshold, while an early-type galaxy is only marginally consistent with the detection limit.  We conclude that FBQ 1633+3134 is a strong candidate for a close separation (0\\farcs66) gravitational lens.  The evidence presented here suggests lensing by a relatively faint ($I>23.0$), $z=0.684$ metal-rich galaxy, although the binary quasar scenario cannot be positively ruled out. Confirmation of the lensing hypothesis can be provided by a deeper radio probe for the fainter component of the system.  We have obtained VLA A array time for late 2000 with this aim in mind, and intend to present these results, along with a systematic analysis of the HST data for this system, in a future paper."
    },
    "0011/astro-ph0011180_arXiv.txt": {
        "abstract": "We investigate the relative sensitivities of several tests for deviations from Gaussianity in the primordial distribution of density perturbations.  We consider models for non-Gaussianity that mimic that which comes from inflation as well as that which comes from topological defects.  The tests we consider involve the cosmic microwave background (CMB), large-scale structure (LSS), high-redshift galaxies, and the abundances and properties of clusters.  We find that the CMB is superior at finding non-Gaussianity in the primordial gravitational potential (as inflation would produce), while observations of high-redshift galaxies are much better suited to find non-Gaussianity that resembles that expected from topological defects. We derive a simple expression that relates the abundance of high-redshift objects in non-Gaussian models to the primordial skewness. ",
        "introduction": "Now that cosmic-microwave-background (CMB) experiments \\cite{boomerang3,maxima,balbi,lange} have verified the inflationary predictions of a flat Universe and structure formation from primordial adiabatic perturbations, we are compelled to test further the predictions of the simplest single-scalar-field slow-roll inflation models and to look for possible deviations.  Measurements of the distribution of primordial density perturbation afford such tests.  If the primordial perturbations are due entirely to quantum fluctuations in the scalar field responsible for inflation (the ``inflaton''), then their distribution should be {\\em very} close to Gaussian (e.g., \\pcite{Guthpi82,Starobinski82,BST83,FRS93,GLMM94,Gangui94,WK00,Ganguimartin99}). However, multiple-scalar-field models of inflation allow for the possibility that a small fraction of primordial perturbations are produced by quantum fluctuations in a second scalar field. If so, the distribution of these perturbations could be non-Gaussian (e.g., \\pcite{allengrinsteinwise87,kofpog88,SBB89,LM97,Pee99a,Pee99b,Salopek99}). Moreover, it is still possible that some component of primordial perturbations are due to topological defects or some other exotic causal mechanism \\cite{bouchet}, and if so, their distribution should be non-Gaussian (e.g., \\pcite{vilke85,vach86,hillscrammfry89,Tur89,albsteb92}).  Detection of any non-Gaussianity would thus be invaluable for appreciating the nature of the ultra-high-energy physics that gave rise to primordial perturbations.  Ruling such exotic possibilities in or out will also be necessary to test the assumptions that underly the new era of precision cosmology.  There are several observables that can be used to look for primordial non-Gaussianity.  CMB maps probe cosmological fluctuations when they were closest to their primordial form, and many authors have developed various mathematical tools to test the Gaussian hypothesis. The statistics of present-day large scale structure (LSS) in the Universe can also be used (e.g., \\pcite{CMLMM93,LS93,Lokas95,Chodbouch96,SP96,Durreretal00,VH00}). The properties and abundances of the most massive and/or highest-redshift objects in the Universe also contain precious information about the nature of the initial conditions (e.g., \\pcite{COS98}; Robinson, Gawiser \\& Silk 1999; Robinson, Gawiser \\& Silk 2000; \\pcite{Willick00,MVJ00} (MVJ); \\pcite{VKMB}). In Verde et al. (2000a; VWHK00), the relative sensitivities of the CMB and LSS to several broad classes of primordial non-Gaussianity were compared, and it was found that forthcoming CMB maps can provide more sensitive probes of primordial non-Gaussianity than galaxy surveys.  Here we extend the results of that paper to include comparisons to the abundances of high-redshift galaxies as well as the abundance and properties of clusters.  One of our original aims was to determine whether any of these probes would be able to detect the miniscule deviations from Gaussianity that arise from quantum fluctuations in the inflaton; unfortunately, we have been unable to find any. Nevertheless, some detectable deviations from Gaussianity are conceivable with multiple-field models of inflation and/or some secondary contribution to primordial perturbations from topological defects. We will follow VWHK00 and parameterize the primordial non-Gaussianity with a parameter that can be dialed from zero (corresponding to the Gaussian case) for two different classes of non-Gaussianity.  We will then compare the smallest value for the parameter that can be detected with each of the different approaches. ",
        "conclusions": "We considered two models for small primordial non-Gaussianity, one in which the primordial density perturbation contains a term that is the square of a Gaussian field (Model A), and one in which the primordial gravitational potential perturbation  contains a term proportional to the square of a Gaussian (Model B). The non-Gaussianity of Model B is precisely that arising in standard slow-roll inflation and in non-standard inflation, while Model A more closely resembles the non-Gaussianity that would be expected from topological defects. We investigated the relative sensitivities of several observables for  testing for deviations from Gaussianity: CMB, LSS and  high-redshift and/or massive objects (e.g., galaxies and clusters).  The analytic tools developed above allow us to address the question of whether the abundance of currently known high-redshift objects can be accommodated within the framework of inflationary models for a given cosmology.  Recently \\scite{Willick00} has studied in detail the mass determination of the cluster MS1054-03 concluding that its mass lies in the range $1.4\\pm 0.3 \\times 10^{15}$ $M_{\\odot}$ for $\\Omega_{\\rm m}=0.3$ (similar to the independent mass estimates by, e.g., \\scite{Tranetal00} and \\scite{Newmanarnaud00}). As already pointed out by \\scite{Willick00}, for $\\Omega_{\\rm m} \\ge 0.3$ the expected number of objects like MS1054-03 in the survey area is $ \\le 0.01$; i.e., it must be a 3-$\\sigma$ fluctuation or larger. Using the formalism we have described here, a primordial non-Gaussianity parameterized by $\\eps_{\\rm B} \\geq 400$ would be required to account for MS1054-03 as a  $1\\sigma$ fluctuation in the $\\Lambda$CDM model described above. This value is much too large  to be consistent with slow-roll inflation. Our calculation shows that if such a non-Gaussianity exists, it would be easily detectable from forthcoming CMB maps."
    },
    "0011/astro-ph0011463_arXiv.txt": {
        "abstract": "{ This paper describes the performance of {\\em XMM-Newton} for serendipitous surveys and summarises the scope and potential of the {\\em XMM-Newton} Serendipitous Survey. The role of the Survey Science Centre (SSC) in the {\\em XMM-Newton} project is outlined. The SSC's follow-up and identification programme for the {\\em XMM-Newton} serendipitous survey is described together with the presentation of some of the first results. ",
        "introduction": "\\label{s1} In contrast with all-sky X-ray surveys which typically provide relatively shallow coverage, serendipitous surveys provide much deeper observations, albeit restricted to smaller sky areas. Serendipitous X-ray sky surveys, taking advantage of the relatively wide field of view afforded by typical X-ray instrumentation, have been pursued with most X-ray astronomy satellites since the {\\em Einstein} Observatory. The resultant serendipitous source catalogues (e.g. EMSS [835 sources] Gioia et al. \\cite{gioia}, Stocke et al., \\cite{stocke}; WGACAT [$\\sim 62000$ unique sources] White, Giommi \\& Angelini, \\cite{WGA}; {\\em ROSAT} 2RXP \\& {\\em ROSAT} 1RXH [$\\sim 95000$ \\& $\\sim 11000$ sources respectively] the {\\em ROSAT} Consortium, 2000) have been the basis for numerous studies and have made a significant contribution to our knowledge of the X-ray sky and our understanding of the nature of the various Galactic and extragalactic source populations.  \\begin{figure*}[htb!] \\parbox{5.7cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1a.eps}}} \\parbox{6.12cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1c.eps}}} \\parbox{5.98cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1e.eps}}} \\parbox{5.67cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1b.eps}}} \\parbox{6.12cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1d.eps}}} \\parbox{5.92cm}{\\resizebox{\\hsize}{!}{\\includegraphics{f1f.eps}}}%  \\caption[] {EPIC images of three fields. Essentially the whole field, 30 arcmin. across, is shown in each case. Top row shows the soft band (0.5-2 keV) images and the bottom row the corresponding hard band (2-10 keV) images. {\\em Left:} Mkn 205, EPIC pn image; {\\em Centre:} OY Car, EPIC MOS1+MOS2 image; {\\em Right:} G21.5$-$0.9 (offset pointing), EPIC pn image. Further details are given in Table~\\ref{tab1}. The images have been smoothed to aid the visibility of faint sources. The trails associated with bright sources in the pn images are due to `out-of-time' events which arrive during the read-out cycle of the CCDs and have not been screened out.} \\label{f1} \\end{figure*}    The {\\em XMM-Newton} Observatory (Jansen et al., \\cite{jansen}), launched in December 1999, provides unrivaled capabilities for serendipitous X-ray surveys by virtue of the large field of view of the X-ray telescopes with the EPIC X-ray cameras (Turner et al., \\cite{turner}; Str\\\"uder et al., \\cite{struder}), and the high throughput afforded by the heavily nested telescope modules. This capability guarantees that each {\\em XMM-Newton} observation provides a significant harvest of serendipitous X-ray sources in addition to data on the original target. This potential is now starting to be realised, as is described below. ",
        "conclusions": ""
    },
    "0011/astro-ph0011149_arXiv.txt": {
        "abstract": "\\noindent Observational data on the large scale structure (LSS) of the Universe are used to establish an upper limit for the amplitude of the tensor mode marginalized over all other cosmological parameters within the class of adiabatic inflationary models. It is shown that the upper 1$\\sigma$ limit for the contribution of a tensor mode to the COBE DMR data is T/S$<1$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/hep-th0011262_arXiv.txt": {
        "abstract": "We critically review several recent approaches to solving the two cosmological constant problems.  The \"old\" problem is the discrepancy between the observed value of $\\rL$ and the large values suggested by particle physics models.  The second problem is the \"time coincidence\" between the epoch of galaxy formation $t_G$ and the epoch of $\\L$-domination $t_\\L$. It is conceivable that the \"old\" problem can be resolved by fundamental physics alone, but we argue that in order to explain the \"time coincidence\" we must account for anthropic selection effects. Our main focus here is on the discrete-$\\L$ models in which $\\L$ can change through nucleation of branes.  We consider the cosmology of this type of models in the context of inflation and discuss the  observational constraints on the model parameters.  The issue of multiple brane nucleation raised by Feng {\\it et. al.} is discussed in some detail. We also review continuous-$\\L$ models in which the role of the cosmological constant is played by a slowly varying potential of a scalar field. We find that both continuous and discrete models can in principle solve both cosmological constant problems, although the required values of the parameters do not appear very natural. M-theory-motivated brane models, in which the brane tension is determined by the brane coupling to the four-form field, do not seem to be viable, except perhaps in a very tight corner of the parameter space. Finally, we point out that the time coincidence can also be explained in models where $\\Lambda$ is fixed, but the primordial density contrast $Q=\\delta\\rho/\\rho$ is treated as a random variable. ",
        "introduction": "The cosmological constant $\\Lambda$ presents us with at least two intriguing problems. Particle physics models suggest that the natural value for this constant is set by the Planck scale, $M_{P}\\sim 10^{18}$ GeV [we use the reduced Planck mass $M_{P}=(8\\pi G)^{-1/2}$]. The corresponding vacuum energy density is $\\rho_\\Lambda\\sim M_{P}^4$, which is some 120 orders of magnitude greater than the observational bounds.  In supersymmetric theories, one can expect a lower value, \\beq \\rho_\\Lambda\\sim\\eta_{SUSY}^4, \\label{susy} \\eeq where $\\eta_{SUSY}$ is the supersymmetry breaking scale. However, with $\\eta_{SUSY}\\gtrsim 1$ TeV, this is still 60 orders of magnitude too high. Until recently, this discrepancy between the expected and observed values was the only cosmological constant problem. Its solution, many believed, was that something so small could only be zero, due to some unknown symmetry or dynamical cancellation.  Thus, it came as a surprise  when recent observations \\cite{Supernova} provided evidence that the universe is accelerating, rather than decelerating, suggesting a non-zero cosmological constant\\footnote{The surprise, however, was not total. In Ref. \\cite{AV95} (well before the supernova data \\cite{Supernova} would give the first observational evidence in this direction) it was noted that anthropic selection effects would place the cosmological constant in the range $\\rho_{\\Lambda}/\\rho_{M0}\\lesssim 10$, and that \"the actual value is likely to be comparable to this upper bound.\" For a flat universe this implies $\\Omega_\\Lambda \\sim 0.9$, not far from the observed value and certainly compatible with it, within the accuracy of the prediction. Similar predictions where made in \\cite{Efstathiou} at about the same time.}. The observationally suggested values of $\\L$ correspond to $\\rho_\\L\\sim\\rho_{M0}$, where $\\rho_{M0}$ is the present density of matter.  This brings yet another puzzle. It is difficult to understand why we happen to live at the epoch when $\\rho_M\\sim\\rho_\\L$.  Another statement of the problem is why the time when $\\Lambda$ starts dominating the universe nearly coincides with the epoch of galaxy formation, \\beq t_\\L\\sim t_G. \\label{tLtG} \\eeq This is the so-called cosmic coincidence problem.  A number of proposed solutions to these problems have recently appeared in the literature \\cite{GV,Bousso,FMSW,Arkani,k-essence,Donoghue}. Some of them rely on some form of the anthropic principle, while others do not.  To our knowlege, the only approach that can explain both puzzles is the one that attributes them to anthropic selection effects.  In this approach, the cosmological constant is assumed to be a random variable that can take different values in defferent parts of the universe.  The purpose of this paper is to give a critical analysis of the proposed approaches, both anthropic and otherwise.  Our main focus will be on the models with a discrete spectrum of $\\L$ which have recently attracted much attention. We shall consider these models in the framework of inflationary cosmology and discuss the calculation of the probability distribution for $\\rL$, as well as the observational constraints on the model parameters.  The paper is organized as follows. In Section II we review the motivation for considering $\\Lambda$ as a random variable. In Section III we discuss models where $\\Lambda$ is a discrete variable, in particular the models where there is a four-form contribution to the cosmological constant, which may relax to a small value through nucleation of branes. In Section IV we analyze the cosmology of such models. In Section V we consider the possibility of coincident brane nucleation. In Section VI we discuss models where the cosmological constant is a continuous variable. In Section VII we consider the possibility of a slowly varying four-form field in theories with extra dimensions. In Section VIII we review some non-anthropic approaches to the problem. In Section IX we consider models where the time coincidence is explained by assuming that the primordial density contrast $Q=\\delta \\rho/\\rho$ (and not necessarily $\\Lambda$) is a random variable. Our conclusions are summarized in Section X. ",
        "conclusions": "The anthropic principle has a bad reputation. It is often regarded as a handwaving argument relying on poorly understood phenomena like intelligent life and having no predictive power. Although this criticism is not entirely ungrounded, there is a class of cosmological models where the use of anthropic principle is not only justified but may in fact be inevitable, and where it can be used to make quantitative predictions. These are the models in which some cosmological parameters, or physical ``constants'', take different values in different parts of the universe.  In such models, one cannot predict the precise values of the parameters that we are going to observe.  One can only hope to calculate the corresponding probability distributions. The criteria for justifying (and compelling) the use of anthropic principle are that the model should provide (i) a mechanism for variation of the parameters and (ii) a way of calculating the probability distributions. Once the probabilities are calculated, one can predict that the parameters are going to be observed within a certain range of values, say, at a 95\\% confidence level.  This seems to be as quantitative as one can possibly get in this class of models.  From a practical point of view, parameters that we can hope to determine anthropically should satisfy the condition that they do not affect life processes, and preferably should also not affect the poorly understood astrophysical processes such as star formation \\cite{VW,GLV,Banks}. So for example the gravitational constant may be hard to determine anthropically with our present level of understanding, since it affects both the evolution of life and star formation. However, gravity does not change chemistry, which is already a big simplification. So it is not inconceivable that the value of Newton's constant may in the future receive an anthropic explanation.  In this paper we discussed anthropic approaches to solving the two cosmological constant problems (CCPs).  The first (old) CCP is the discrepancy between the observed small value of $\\rL$ and the large values suggested by particle physics models.  The second (time coincidence) CCP is the puzzling coincidence between the epoch of galaxy formation $t_G$ and the epoch of $\\L$-domination $t_\\L$.  While it is conceivable that the old CCP can be resolved by fundamental physics \\cite{Witten,Dvali,Arkani,Guendelman}, we have argued that the time coincidence problem calls for an anthropic explanation.  We first considered models with a discrete spectrum of $\\rL$. These include Abbott's scalar field model with a \"washboard\" potential \\cite{Abbott}, as well as models in which $\\rL$ can change through brane nucleation accompanied by a change of the four-form field $F$ \\cite{Teitelboim,Bousso,FMSW,Donoghue}. Such models can solve both CCPs, provided that {\\it (i)} the separation between the discrete values of $\\rL$ is $\\epsilon\\lesssim \\rho_{M0}$, where $\\rho_{M0}$ is the present matter density, {\\it (ii)} the probability distribution for $\\rL$ at the end of inflation is nearly flat, ${\\cal P}_*(\\rL)\\approx {\\rm const}$, and {\\it (iii)} the brane nucleation rate is sufficiently low, so that the present vacuum energy does not drop significantly in less than a Hubble time.  We discussed the cosmology of this class of models, the calculation of the prior distribution ${\\cal P}_*(\\rL)$, and the observational constraints on the model parameters.  The required values of the \"level separation\" $\\epsilon$ may appear uncomfortably small, but Feng, March-Russell, Sethi and Wilczek (FMSW) \\cite{FMSW} have argued that they can naturally arise due to non-perturbative effects in M-theory. In M-theory-related models, the brane tension $\\sigma$ is related to $\\epsilon$ through $\\epsilon\\sim \\sigma\\rho_{bare}^{1/2}/M_p$ and should also be very small. Our analysis shows that in such models the conditions {\\it (i)-(iii)} cannot be satisfied without fine-tuning of the parameters.  It was conjectured by Weinberg \\cite{Weinberg87} that the condition {\\it (ii)} of a flat {\\em a priori} distribution for $\\Lambda$ would automatically be satisfied in any particle physics model where the cosmological constant is a random variable. In Ref. \\cite{GV} we showed that this conjecture is not always satisfied in models where the role of the cosmological constant is played by a slowly varying field. Here, we have shown that the conjecture is generically not satisfied in four-form models either. In fact, this condition has to be enforced in order to fit observations. This, in turn, places severe constraints on the model parameters. Hence, in trying to solve the cosmological constant problems by anthropic means, the flat {\\it a priori} distribution for $\\L$ cannot be taken for granted and the problem of calculating ${\\cal P_*}$ has to be addressed.  Bousso and Polchinsky \\cite{Bousso} have studied models with multiple four-form fields $F_i$ and found that the spectrum of the allowed values of $\\rL$ can be sufficiently dense even for large brane tensions. However, in this case the vacua with nearby values of $\\rL$ have very different values of $F_i$, and a flat probability distribution required in {\\it (ii)} is rather unlikely.  Moreover, the low-energy physics in such vacua is likely to be different, and it appears that the anthropic approach to solving the CCPs cannot be applied in this case \\cite{Banks}.  For models unrelated to M-theory, $\\sigma$ and $\\epsilon$ are generally unrelated, and values consistent with the constraints {\\it (i)-(iii)} can easily be found.  However, if one gives up the M-theory connection, then the FMSW argument cannot be used, and one has to seek an alternative explanation for the tiny value of $\\epsilon$.  Alternatively, one might seek modifications of the FMSW model that could relax the relation between $\\epsilon$ and $\\sigma$.  All of the earlier discussions of the cosmology of discrete $\\Lambda$ models encountered the \"empty universe problem\" \\cite{Banks1,Abbott,Teitelboim,Bousso,FMSW,Donoghue,Banks}. In order to make the present vacuum sufficiently stable, the brane nucleation has to be strongly suppressed.  One then finds that the time it takes the universe to evolve from some initial high value of $\\rL$ to the present low value is much greater than the present Hubble time.  This suggests that by the time the process is complete, any matter that the universe initially had may get diluted to an extremely low density, so that one would end up with an empty universe dominated by the cosmological constant.  We have argued that the empty universe problem disappears when the eternal nature of inflation is taken into account.  During inflation, brane nucleations leading to higher and lower values of $\\rL$ have nearly equal probabilities.  As a result, the values of $\\rL$ are randomized, with different parts of the universe thermalizing with different values.  The resulting probability distribution ${\\cal P}_*(\\rL)$ can be calculated using the stochastic formalism we developed in Section IV.  The slow rate of brane nucleation is not a problem in eternal inflation, since an unlimited amount of time is available.  FMSW suggested an interesting possibility that nucleation of multiple branes could be enhanced by a large degeneracy factor due to the light fields living on the branes.  If true, this could significantly modify the brane model cosmology.  In Section V we studied multiple brane nucleation in some detail and found that the pre-exponential factor in the brane nucleation rate can both enhance and suppress multiple brane nucleation, depending on the field content of the branes.  We also concluded that models in which multiple brane nucleation dominates can be ruled out, because in such models there is nothing to prevent our present vacuum from tunneling down to deep anti-de Sitter space.  We also discussed models with a continuous spectrum of $\\rL$, in which the role of the cosmological constant is played by the potential $V(\\phi)$ of a scalar field $\\phi(x)$.  The potential has to be very flat, so that its value does not significantly evolve on the present Hubble time scale.  The values of the field $\\phi$ are randomized by quantum fluctuations during inflation, and models can easily be constructed in which the resulting probability distribution for $V(\\phi)$ is nearly flat in the range of interest, thus solving both CCPs.  The challenge here is to justify the very flat potentials required in this class of models.  Possibilities include a pseudo-Goldstone field which acquires a potential through instanton effects \\cite{GV}, a large running of the field renormalization \\cite{Weinbergcomment}, and a non-minimal kinetic term with an exponential $\\phi$-dependence \\cite{Donoghue}.  We have pointed out some difficulties of the instanton approach.  We thus see that both discrete and continuous $\\Lambda$ models could in principle solve both of the CCPs.  However, none of the models that have been suggested so far appears particularly well motivated or natural.  An alternative approach is to assume that the old CCP can be solved within the fundamental theory.  The cosmological constant is then truly a constant and is given by an expression like $\\rL\\sim M_w^4 M_P^{-2}$, as in \\cite{Arkani}.  At the same time, the amplitude of density fluctuations $Q$ could be a random variable, so that the epoch of galaxy formation $t_G$ is different in different parts of the universe. We have shown in Section VIII that, for a wide class of prior distributions ${\\cal P}_*(Q)$, most of the galaxies will be in regions where $t_G\\sim t_\\L$, thus explaining the cosmic time coincidence. It would be interesting to extend this analysis and calculate the distribution ${\\cal P}_*(Q)$ for some models with a variable $Q$.  One would then have some idea of how naturally the distributions of the required type can be obtained."
    },
    "0011/astro-ph0011419_arXiv.txt": {
        "abstract": "{ In this paper, we present our preliminary results on Seyfert galaxies with circumnuclear/nuclear starburst(SB) activity. We have searched the recent available literature and found 76 active galaxies with clear evidence of nuclear SB activity, among which 16 are Seyfert 1s, 51 Seyfert 2s, and 9 LINERs. After studying the 51 Seyfert 2s, we find that those Seyfert 2s with hidden Seyfert 1 nuclei, have similar Infrared-Radio properties as Seyfert 1 galaxies, and are different from \"real\" Seyfert 2s without a hidden Seyfert 1 nucleus. The later are similar to starburst galaxies. } ",
        "introduction": "The connection between starburst(SB) activity and AGNs is one of the most important and hotly debated issues in the study of active galaxies. On the theoretical side, Norman \\& Scoville (1988) have suggested that the circumnuclear star-forming activity could drive gas into the innermost nuclear region to feed the central black hole. However, Terlevich (1992) and Terlevich \\etal \\ (1992) have proposed that Seyfert 2s and LINERs could be well reproduced by violent star formation in a metal-rich environment, such as in the nuclei of early-type galaxies, which has been well known as the Starburst model for AGN (see e.g. the recent review by Cid Fernandes 1997 and references therein).  In a series of papers(Dultzin-Hacyan  1995, Dultzin-Hacyan \\& Ruano  1996, and Dultzin-Hacyan \\etal \\  1999), it has been shown that there are several lines of evidence against the simplest formulation of a \"Unified Scheme\" for Seyfert galaxies, according to which, {\\it all} Seyfert 2s have a hidden Sy 1 nucleus. In this paper, we add new evidence to the existence of two kinds of Sy 2s: The hidden Sy 1, and the \"real\" or \"pure\" Sy 2. According to these authors the difference between Sy 1 and \"real\" Sy 2 is relative decrease in the accretion power in the first versus a relative increase of nuclear/circumnuclear SB activity in the second.  On the other hand, more and more high-quality observations from HST and large telescopes on ground, support the presence of active star formation activity around the nuclei in dozens of Seyfert galaxies. For example, in NGC 1068, a prototype of Seyfert 2 galaxy, the circumnuclear starburst accounts for at least 50\\% of the whole bolometric luminosity, 81\\% at far UV($\\sim$ 1500 \\AA), and 83 \\% at near UV ($\\sim$ 2500 \\AA) bands(Gonzalez Delgado  \\etal \\ 1997). And Heckman \\etal \\ (1997) and Gonzalez Delgado \\etal \\ (1998) have presented high resolution UV images of 4 Seyfert 2 galaxies, and found compact nuclear starburst in all of them.  In Seyfert 2 galaxies with nuclear starburst, the observed UV fluxes, as suggested by Colina \\etal \\ (1997), are dominated not by AGN, but by the young massive stars in the nuclear region.  There are several ways to detect star-forming activities in Seyfert galaxies (Wilson  1987). Narrow-band \\Ha \\ imaging (Evans \\etal \\  1996, Gonzalez Delgado \\etal \\  1997), high-resolution radio continuum mapping(Forbes \\& Norris  1998), high S/N spectroscopy on central region and using stellar population synthesis technique (Schmitt \\etal \\  1999, Boisson \\etal \\  2000, and Jimenez-Benito \\etal \\  2000), the shape of the IRAS spectra, and infrared(IR) spectroscopics searching for absorption lines from red supergiants (Terlevich \\etal \\ 1990, Oliva \\etal \\ 1995).  In order to study the role of starburst in Seyfert galaxies and the connection between starburst and AGNs, we collect all active galaxies with circumnuclear and/or nuclear starburst as completely as possible from the recent literature, which are listed in Table 1.  This paper is organized as follows. In Section 2, we present the basic information of our sample of 76 active galaxies with circumnuclear/nuclear starburst, and the IR and radio properties of 51 Seyfert 2 galaxies are presented and discussed in Section 3. Finally, Section 4 summarizes our main results. ",
        "conclusions": "We have collected 76 active galaxies with circumnuclear or nuclear star-forming activities, among which, 16 are Seyfert 1s, 51 Seyfert 2s and 9 LINERs. After studying their IR and radio properties of 51 Seyfert 2s, we confirm that there are two distinct types of Seyfert 2 galaxies, those with hidden Seyfert 1 nuclei share similar properties with Seyfert 1s, and those without, are similar to Starburst galaxies."
    },
    "0011/astro-ph0011133_arXiv.txt": {
        "abstract": "We investigate the consequences of a continuously injecting central engine on the gamma-ray burst afterglow emission, focusing more specifically on a highly-magnetized millisecond pulsar engine.  For initial pulsar parameters within a certain region of the parameter space, the afterglow lightcurves are predicted to show a distinctive achromatic bump feature, the onset and duration of which range from minutes to months, depending on the pulsar and the fireball parameters. The detection of or upper limits on such features would provide constraints on the burst progenitor and on magnetar-like central engine models. An achromatic bump such as that in GRB 000301C afterglow may be caused by a millisecond pulsar with $P_0=3.4$ms and $B_p=2.7\\times 10^{14}$ G. ",
        "introduction": "Much of the current research on Gamma-Ray Bursts (GRB) is aimed at determining the nature of the central engine and its progenitor system. While recently substantial results have begun to accumulate, the evidence is still tentative.  Thus, the investigation of criteria to differentiate between various central engine possibilities is desirable.  Almost all the present fireball models, including those considering various non-uniform injection scenarios, assume that the energy injection into the fireball occurs in a short period of time. This is also the case in the ``refreshed shock'' scenario (Rees \\& M\\'esz\\'aros 1998; Kumar \\& Piran 2000; Sari \\& \\Mesz~ 2000). However, in some types of central engines, such as a fast-rotating high-field pulsar (magnetar) or a black hole plus a long-lived debris torus system, a significant energy input into the fireball may in principle continue for a time scale significantly longer than the $\\gamma$-ray emission. Therefore, there is a need to investigate a continuously-fed fireball in more detail.  An additional motivation is provided by the recent detection of Fe features in the X-ray afterglow of GRB 991216 after about 1.5 days (Piro et al.  2000) and GRB 000214 after about 1 day (Antonelli et al, 2000), which may require a continuing post-burst outflow in order to achieve less restrictive Fe abundance constraints (Rees \\& \\Mesz, 2000).  Dai \\& Lu (1998a,b) first considered continuous injection from a millisecond pulsar to interpret the afterglow light curves of some GRBs, but did not perform a systematic study on this topic.  In this paper, we investigate the observational consequences of a continuously injecting central engine, and more specifically, focus on the possibility that the central engine is a millisecond pulsar, in particular a highly magnetized pulsar or magnetar. ",
        "conclusions": "If the central engine of GRB are fast-rotating pulsars, the afterglow lightcurves may show a distinctive, achromatic feature, for pulsars whose initial parameters are within ``pulsar signature region'' region defined by equations (\\ref{1}), (\\ref{2}) and (\\ref{3}). In this case the afterglow light curves flatten after a critical time $T_c$ (eq.[\\ref{tec}]), and steepen again after a time ${\\cal T}_{em}$ (eq.[\\ref{Teem}]).  This region of pulsar parameter phase space includes ``magnetars\" or ultra-high field pulsars, which play a large role in some GRB pulsar progenitor models.  The current data on early afterglows are insufficient to provide good tests for this feature. However, an interesting possibility is the achromatic bump observed in the afterglow of GRB 000301C, which for a limited time deviates significantly from the standard broken power-law fit (Masetti et al. 2000).  Possible explanations include running into a non-uniform ambient density (Berger et al. 2000), and modification by a microlensing event (Garnavich, Loeb \\& Stanek 2000). We suggest here a third possibility, that the bump may be caused by the pulsar signature discussed above. Taking (Berger et al 2000) $\\alpha_1 \\sim-1.28$ for the principal temporal index (before the bump and before the decay ascribed to a jet transition), the temporal index during the pulsar signature is expected to be $\\alpha_2\\sim 0.15$, which seems reasonable to fit the achromatic bump. Since the feature occurs a couple of days after the GRB trigger, the initial pulsar parameters are in regime II. Taking $E_{\\rm imp,51}\\sim 1.1$ from the observations (Berger et al. 2000), the pulsar parameters follow from setting ${\\cal T}_{em}\\sim3.8$ d and $T_c\\sim2.5$ d, which gives $P_0\\sim 3.4$ ms and $B_{p,15}\\sim 0.27$ (Fig.1). Detailed $\\chi^2$ fit to the data would be necessary to validate this proposal.  Another relevant observation may be the recent Fe line detection in the X-ray afterglow of GRB 991216 (Piro et al 2000), which in one interpretation would require a continuously-injecting central engine (Rees \\& M\\'esz\\'aros 2000). If $L_{em}\\sim 10^{47} {\\rm erg~s^ {-1}}$ is assumed 37 hours after the burst (Rees \\& M\\'esz\\'aros 2000), this continued injection could be due to a pulsar with $B_{p,15}\\sim 0.15$ and $P_{0}<1.2$ ms, which would imply a signature bump at $\\simg 1$ hour.  Such a bump is not seen in this afterglow (e.g. Halpern et al. 2000).  However, a slightly weaker luminosity (e.g. $L_{em} \\leq 3\\times 10^{46} {\\rm erg~s^ {-1}}$) could also explain the Fe line features, by assuming a slightly larger Fe abundance (which is very low in this model). The required pulsar can therefore be more magnetized, and the characteristic times for the signature bump would be expected early enough to have evaded detection in this GRB.  In conclusion, the detection of or upper limits on such characteristic afterglow bumps by missions such as {\\em HETE2} or {\\em Swift} may be able to provide interesting constraints on magnetar GRB models and their progenitors."
    },
    "0011/astro-ph0011305_arXiv.txt": {
        "abstract": "We present the first long-term optical and infrared study of the optical counterpart to the source RX J0812.4--3114, an X-ray pulsar with a Be type companion. During the period covered by the observations the profile of some Balmer lines changed from absorption to emission and back again to absorption. Contemporaneously, the infrared magnitudes varied by more than 0.8 mag. This long-term variability is interpreted as the formation and subsequent dissipation of the Be star's disc. The building up of the disc ended up in an active X-ray state characterised by regular outbursts occurring at 80 day intervals. The overall duration of the formation/dissipation of the disc is found to be $\\simless 4.3 $ years. Optical spectroscopic and infrared photometric observations were used to refine the spectral type of the primary (B0.2IV) and to monitor the circumstellar  envelope around the Be star. UBVRI and $uvby\\beta$ photometric observations allowed the determination of the astrophysical parameters of the optical companion. ",
        "introduction": "The star LS 992 is the optical counterpart to the X-ray source RX J0812.4--3114. Such an association was originally proposed by Motch et al. (1997), as a result of a systematic cross-correlation between the ROSAT All Sky Survey (Voges et al. 1999) and several OB star catalogues in the SIMBAD database.  The X-ray light curve of \\ls\\ is characterised by 31.88 second pulsations, while the X-ray spectrum is best represented by an absorbed power-law component with a exponentially cut-off (Reig \\& Roche 1999). In December 1997 the source made a transition from a quiescent state to a flaring state (Corbet \\& Peele 2000), in which regular flares separated by 80 day intervals were detected with the All-Sky Monitor (ASM) onboard the {\\em Rossi X-ray Timing Explorer}. Corbet \\& Peele (2000) attributed the origin of these flares to the periastron passage of the neutron star, hence this periodicity was naturally associated with the orbital period. The peak X-ray luminosity during outbursts was estimated to be of the order of 2.3 $\\times$ 10$^{36}$ erg s$^{-1}$, assuming a distance to the source of 9 kpc.  In the optical band, LS~992 appears as a relatively bright $V\\sim$~12.4 B0--1III-V star (Reed 1990). The optical spectra revealed H$\\alpha$ in emission and evidence for short-term ($\\sim$ weeks) V/R variability (Motch et al. 1997).  The combination of an emission line B star plus a neutron star, whose existence is inferred from the characteristics of the X-ray emission, defines a Be/X-ray binary. Be/X-ray  binaries  comprise approximately  70\\% of the more  general class of high mass X-ray binaries,  the remaining  $\\sim$  30\\%  contain evolved  (luminosity  class I and II) primaries. The high energy radiation is the result of accretion of the material  expelled by the Be star onto the neutron star.  In this work we present the first long-term study of the optical counterpart of RX J0812.4--3114 and discuss the consequences of the correlated behaviour in the optical, infrared and X-ray wavelength bands. ",
        "conclusions": "We have witnessed the appearance and disappearance of the circumstellar envelope in the Be/X ray binary \\ls.  Using $UBVRI$ and $uvby\\beta$ photometric observations we have derived the astrophysical parameters of the Be/X-ray binary \\ls. The optical companion turns out to be a B0.2IVe star showing marked optical and infrared variability on timescales of months to years. The reversal from positive to negative and then back to positive of the equivalent width of the H$\\beta$ line, the large amplitude of variation of both, the H$\\alpha$ equivalent width ($\\Delta EW(H\\alpha) \\approx 20$ \\AA) and the infrared magnitudes ($\\Delta J \\approx \\Delta H \\approx \\Delta K \\approx 0.8$ mag) are interpreted as the formation and dissipation of the Be star's disc. The long-term monitoring of \\ls\\ constrains the timescale for formation and dissipation of the disc to about 4.3 years. Measurements found in the literature seem to indicate that this episode is cyclic and basically supports the picture of dynamically unstable discs, that is, the discs go through cycles of dissipation/reformation which control the X-ray behaviour of the source."
    },
    "0011/astro-ph0011075_arXiv.txt": {
        "abstract": "The present study deals with the problem of deriving the coefficients of the fundamental plane (FP) of early-type galaxies.  We introduce statistical models of the FP and relative fitting methods: the MIST ({\\it Measurement errors and Intrinsic Scatter Three dimensional}) fits\\footnote{The FORTRAN codes of the MIST algorithms are available on request to the authors.}. The MIST fits account for the measurement errors on the variables and their correlations as well as for the intrinsic scatter. We show that the lack of a model of the intrinsic scatter of the FP is the origin of the systematic differences between the various fitting methods.  We also address the problem of estimating the uncertainties of the FP coefficients and determine a simple relation between the sample size and the expected accuracy of the coefficients.  The present study leads to define a `minimum sample size' for a correct estimate of the uncertainties. For $N \\widetilde{<}30$, both theoretical formulae and re-sampling techniques, like the bootstrap, do not give reliable estimates.  The question of the 'universality' of the FP is addressed by applying the MIST fits to ten samples of cluster galaxies. The FP slopes are actually consistent between the different samples, but, due to the large uncertainties, they could also hide significant systematic differences.  The feasibility of the measurement of the possible variations of the FP slopes as a function of redshift is also proved. ",
        "introduction": "The fundamental plane (FP) is a bivariate relation between observed global properties of early-type (E) galaxies like the effective radius $r_e$, the mean surface brightness $\\mu_e$, and the central velocity dispersion $\\sigma_0$ (e.g. Dressler et al. \\cite{dlb87}, Djorgovski \\& Davis \\cite{djd87}) of the form\\footnote{This representation of the FP is adopted throughout the paper.}: \\begin{equation} \\log r_\\mathrm{e} = a \\log \\sigma_0 + b \\mu_\\mathrm{e} + c\\ . \\label{fpeq} \\end{equation} The main features of the FP are its small scatter ($\\sim 10\\%$ in $r_e$, of the same order of the measurement errors) and the `tilt' of its slopes with respect to the prediction of the virial theorem for a family of homologous systems with constant $M/L$ ratio.  It is generally believed that the `tilt' carries information on the nature of Es, although no convincing interpretation has been found that can be reconciled with the small scatter (see e.g. Ciotti et al. \\cite{clr96}, Renzini \\& Ciotti \\cite{rec93}). It is now clear that part of the `tilt' is due to the intrinsic non-homology, both structural and dynamical, of the E family (see e.g. Busarello et al. \\cite{bcc97}, Graham \\& Colless \\cite{grc96}, Prugniel \\& Simien \\cite{prs97} and references therein). Stellar population effects (e.g. systematic differences in age/metallicity and interstellar matter content) should account for the remaining tilt (see e.g. Mobasher et al. \\cite{mga99}). In that case, the tilt should decrease when moving to near-infrared wavelengths, where these effects are minimized: some evidences exist that this is actually the case (e.g. Pahre et al. \\cite{pcd98}, Scodeggio et al. \\cite{sgb98} and references therein). The wavelength dependence of the FP enters also in the comparisons between FP determinations at different redshifts, where different rest-frame wavebands are sampled.  The FP has been soon recognized as a precise tool for distance determinations (e.g. Dressler et al. \\cite{dlb87}, Hudson et al. \\cite{hls97}) and, more recently, for constraining cosmological parameters and for studying galaxy evolution (van Dokkum et al. \\cite{vdf96}, Bender et al. \\cite{bsz98}, J\\o rgensen et al. \\cite{jor99}, Kelson et al. \\cite{kivd00}, and references therein).  All the above applications are based on comparisons between different determinations of the FP: e.g. different samples, different wavebands, and different redshifts. The method used to derive the FP coefficients and the relative uncertainties play therefore a central r\\^ole for a proper use of the FP.  The determination of the FP requires considerable observational and data analysis efforts (see e.g. Ziegler \\& Bender \\cite{zib97} and Ziegler et al. \\cite{zsb99}), it is thus desirable that a similar effort be made for the derivation of the coefficients of the FP and for an accurate estimate of the relative uncertainties.  There is still no agreement in the scientific community about the fitting method to adopt. This is partly due to the numerous applications for which the FP is derived (see above), but it also originates from two important points: (1) the measurement errors on the FP variables are comparable (see e.g. Smith et al. \\cite{slh97}, hereafter SLH97) and correlated (see J\\o rgensen et al. \\cite{jfk95a}, hereafter JFK95a), (2) the scatter around the plane is not completely accounted for by the measurement uncertainties but has also an intrinsic origin (see e.g. J\\o rgensen et al. \\cite{jfk96}, hereafter JFK96).  In  the present work we address the problems relative to: (I) the choice of the procedure to derive the FP coefficients (Sect.~2) and (II) techniques to estimate the corresponding uncertainties (Sect.~3).  Starting from the bi-dimensional models introduced by Akritas \\& Bershady (\\cite{akb96}, hereafter AkB96), we propose in Sect.~2.1 statistical models of the FP that account for the various sources of scatter on the variables and derive the relative fitting procedures: the MIST ({\\it Measurement errors and Intrinsic Scatter Three dimensional}) fits.  We then discuss in Sect.~2.2 the capability of the fitting procedures to derive the coefficients of the FP.  Section 3 deals with the uncertainties of the coefficients of the FP. The simulation algorithm used for the present analysis is described in Sect.~3.1. In Sect.~3.2 the different methods to estimate the uncertainties are analyzed and compared to the results of the simulations.  The results and the techniques developed in Sect.~2, 3, are applied in Sect.~4 to the comparison of the FPs of different clusters of galaxies. ",
        "conclusions": "The FP of cluster galaxies is one of the most promising tools to study galaxy evolution, to constrain the epoch of galaxy formation, and to set constraints on the geometry of the universe. To those aims the FP has been recently studied in intermediate-redshift ($z=0.18 - 0.83$) clusters (van Dokkum \\& Franx \\cite{vdf96}, Kelson et al. \\cite{kivd97}, Bender et al. \\cite{bsz98}, van Dokkum et al. \\cite{vdf98}, J\\o rgensen et al. \\cite{jor99}, Kelson et al. \\cite{kivd00}). The main step in the above applications is the comparison between FP determinations in different clusters, wavebands and redshifts, for which the determination of the FP coefficients and of the relative uncertainties plays a crucial r\\^ole.  The present study aims at clarifying various aspects in the determination of the FP coefficients and uncertainties.  \\subsection{The problem of fitting the FP.}  We introduced a statistical model of the FP that takes into account the measurement errors on the variables and the intrinsic scatter. The correlation between measurement uncertainties is also accounted for.  We derived the MIST$_i$ (\\emph{Measurement errors and Intrinsic Scatter Three dimensional}) fitting procedures, where $i = \\log r_{\\mathrm{e}}, \\mu_{\\mathrm{e}}, \\log \\sigma_0$ is the dependent variable of the fit. A bisector method, \\erbi, is also introduced.  The assumptions under which the MIST procedures give unbiased estimates of the FP coefficients are that the errors and the intrinsic scatter do not depend on the `location' on the fitting plane (H1), their  covariance matrices are known (H2 and H3), and the covariance matrices of the observed quantities are known (H4).  Under these assumptions, the fits give the same coefficients whatever the choice of the dependent variable.  For what concerns the above assumptions, the knowledge of the average measurement errors can be easily introduced in the procedures in order to satisfy H1 and H2, while the problems posed by H4 can be overcome by assuming a normal distribution for the observed quantities.  Assumption H3, on the other hand, cannot be satisfied since we do not have a model for the intrinsic dispersion of the FP variables. We showed that it is precisely the lack of such a model that biases the best-fit coefficients in such a way that their values depend on the choice of the dependent variable and could not coincide with the slopes of the `true relation'.  \\subsection{Uncertainties on the FP coefficients.}  We addressed the problem of the estimate of the uncertainties on the FP coefficients by numerical simulations based on the best sampled FP so far (the Gunn $r$ FP of Coma by JFK96).  The results of the simulations are presented in Fig. 4 in terms of $\\delta a/a$ and $\\delta b/b$ versus $ \\log( N )$. This figure can be used to state the sample size needed to achieve a given accuracy for the FP coefficients.  On the basis of the simulations we tested the performances of the methods used to estimate the uncertainties: the theoretical formulae and the bootstrap technique.  We showed that for $N > N_{min} \\approx 30$ the theoretical formulae can recover the variances of the FP coefficients within about 5\\%, while for $N < N_{min}$ the uncertainties are systematic smaller than the true values and, what is more troublesome, are affected by a large scatter.  The bootstrap technique gives, on the average, more accurate estimates than theoretical methods but for small samples the scatter of the individual values is much larger. Again, a minimum sample size of $N = N_{min} \\approx 30$ exists for the bootstrap to give reliable estimates.  \\subsection{Implications for the use of the FP.}  As an application of the present study we addressed in Sect. 4 the question of the `universality' of the FP. The main obstacles are the limited size of the available (homogeneous) samples and their different completeness characteristics.  We find that the slopes of the FP in the ten clusters considered are consistent \\emph{ within the large uncertainties} (due to the small number of galaxies in most of the samples). We stress that this result should be taken at most as an indication in favor of the universality of the FP and not as an evidence.  With the available data samples it is not possible to settle the question of the universality: to this aim, larger and more homogeneous samples would be needed.  For the application of the FP to the study of galaxy evolution it would be crucial to know the behavior of the slopes as a function of redshift. Pahre et al. (1998) built a model for the FP that predicts a change in slope with redshift: if age contributes to the tilt of the FP, massive galaxies should be much older than the least luminous ones. As a consequence of this `differential' evolution, the slopes will change with redshift. They predict the change with redshift of the coefficient $a$ of the FP (see their Fig. 10) by adopting two different models for the stellar evolution (Bruzual \\& Charlot \\cite{brc98} and Vazdekis et al. \\cite{vcp96}). The Bruzual \\& Charlot models produces a rapid change of $a$, that should decrease by about 0.15 from z=0 to z$\\sim$0.3, namely, 12\\% of its local value in the Gunn $\\mathrm{r}$ band.  Recently, Kelson et al. (\\cite{kivd00}) derived the FP for a cluster at z=0.33 with 30 galaxies. The authors claim that the slopes are fully consistent with those of the local FP (Coma). They derive the FP by the orthogonal fit and give the following uncertainties for the slopes: $\\delta a/a$=0.1 and $\\delta b/b$=0.11, that are consistent with the values obtained here from the simulations with $n$=30 (see Fig. 4). These data seem thus to rule out the rapid evolution of $a$ predicted with the Bruzual \\& Charlot model.  The situation is quite different with the Vazdekis et al. models. At $z=0.3$ the slope $a$ decreases by only 0.03\\%. The minimum sample size needed to achieve such a precision at $z=0.3$ would be $N \\approx 300$, which is out of the reach of any possible observation. The situation improves by moving at higher redshifts, where the magnitude of the predicted change of the slopes is larger. For the cluster at z=0.83 studied by van Dokkum et al. (\\cite{vdf98}), in fact, the situation is by far more promising. To test the predicted change in slope it would be sufficient to determine the FP on $N \\approx 40$ galaxies, that is well within the possibility of a 8-10m class telescope in a reasonable number of nights."
    },
    "0011/astro-ph0011243_arXiv.txt": {
        "abstract": "We propose a cosmological model that makes a significant step toward solving the coincidence problem of the near similarity at the present of the dark energy and dark matter components. Our cosmology has the following properties: a) among flat and homogeneous spaces, the present universe is a global attractor: {\\it all} the possible initial conditions lead to the observed proportion of dark energy and dark matter; once reached, it remains fixed forever; b) the expansion is accelerated at the present; c) the model is consistent with the large-scale structure and microwave background data; d) the dark energy and the dark matter densities  scale similarly after equivalence and are close to within two orders of magnitude. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011486_arXiv.txt": {
        "abstract": "We present the initial results of a study aimed at computing the global effects of the X-ray ionization due to young stars within a molecular cloud that forms stars with a standard IMF at a rate accelerating with time. We model the gas distribution of the cloud as a fractal of dimension 2.3. We introduce the concept of a {\\it R\\\"ontgen sphere} as the region around each PMS star where the ionization rate due to stellar X-rays exceeds that due to interstellar cosmic rays.  Using values of the X-ray luminosity typical of T Tauri stars, we find that within the R\\\"ontgen radius the average ionization rate is increased by a factor 4 to 20 with respect to the value obtained from cosmic-rays. ",
        "introduction": "The initial results of the {\\em Einstein Observatory} have shown that young stars are powerful X-ray emitters (Montmerle et al. 1983). It was then noted that the incident X-ray flux from YSOs may be the dominant controlling agent of the cloud ionization (Krolik \\& Kallman 1983). The possibility that X-ray ionization constitutes a natural feedback mechanism in the physics of star formation was suggested by Silk \\& Norman (1983). They argued that, in the framework of magnetically controlled cloud evolution and collapse, the rate of star formation both determines, and is determined by, the cloud ionization level. On a smaller scale, also the chemistry of the cloud around the immediate environs of a low-mass young star is deeply affected by the interaction of the high-energy photons with the dense molecular gas (Glassgold, Najita \\& Igea 1997).  The aim of this work is to re-examine the feedback mechanism of stellar X-rays on ambient gas in light of current views of the history of star formation in clusters and associations (e.g. Palla \\& Stahler 2000) and of the expanded knowledge of the X-ray properties of both embedded and optically revealed young stars (Feigelson \\& Montmerle 1999).  Observational studies in a variety of star forming regions  have indicated that the X-ray activity lasts longer and at higher levels than estimated in the initial studies. Also, it appears that molecular clouds do not spawn stars at constant rates during their lifetimes, but with a seemingly accelerating pattern. Finally, observations of molecular clouds have shown that with increasing spatial resolution the gas distribution breaks up into substructure of yet smaller scales.  This suggests that the density distribution may have a  fractal structure (e.g. Falgarone, Phillips \\& Walker 1991).  Obviously, all these properties can enhance the impact of the interaction of the stellar X-rays with the surrounding medium. To assess such effects quantitatively, we have first examined the distribution of X-ray energy around a star with a characteristic X-ray luminosity $L_{\\rm X}/L_{\\rm bol} \\la 10^{-3}$, typical of T-Tauri stars.  We have then computed the overall contribution of the X-ray ionization due to an ensemble of stars formed inside a molecular cloud of sizes between 0.5 and 2 pc.  The basic question we want to address is: {\\it what fraction of the mass of the cloud is affected by the enhancement in the ionization level due to X-rays?}  \\section {The R\\\"ontgen Radius} The ionization rate, $\\zeta_{\\rm X}$, due to X-rays from a YSO as a function of  distance, $r$, from the star can be evaluated from the expression:  \\begin{equation} \\zeta_{\\rm X} \\ =\\  1.7 \\,\\frac{L_{\\rm X} \\tilde{\\sigma}}{4\\pi r^2 \\Delta \\epsilon} \\ \\frac{\\int_{\\nu_0}^{\\infty} J_{\\nu}\\, e^{- \\tau_{X}(\\frac{\\nu}{\\nu_{\\rm X}} )^{-n}} \\big{(} \\frac{\\nu}{\\nu_{\\rm X}} \\big{)}^{-n}  d \\nu}{\\int_{\\nu_0}^{\\infty} J_{\\nu} \\, d\\nu} \\,\\ {\\rm s}^{-1} \\end{equation}  \\noindent where $L_{\\rm X}$ is the X-ray luminosity of the star, $\\sigma(E) =\\tilde{\\sigma}(E/1 {\\rm keV})^{-n}$ is the total photoelectric cross section with $n= 2.51$ and $\\tilde{\\sigma}=2.16\\times 10^{-22}$ cm$^2$. $\\Delta \\epsilon=35$ eV is the mean energy to make a ion pair in weakly-ionized gas with cosmic abundance, and the factor 1.7 accounts for the ionization due to secondary electrons (Cravens \\& Dalgarno 1978). In eq. (1), $\\tau_{\\rm X}=n_{\\rm H}\\tilde{\\sigma}r$ is the optical depth at $h\\nu$=1 keV, $J_{\\nu}$ is the monochromatic flux   emitted from the star, in the standard assumption of a thermal bremsstrahlung spectrum. The minimum photon energy in the integral of eq. (1) is $h\\nu_0=$ 0.1 keV. Assuming an X-ray luminosity of L$_{\\rm X}=10^{30}$ erg s$^{-1}$, the resulting run of the ionization rate with distance from the central star is displayed in Figure 1.  \\begin{figure}[h] \\plotfiddle{LORENZANIA1.eps}{6cm}{0}{65}{55}{-160}{-7} \\caption{The X-ray ionization rate versus distance from the star, for a YSO with $L_{\\rm X}=10^{30}$ erg s$^{-1}$ and a two temperature spectrum model. $R_{\\rm X}$ is the R\\\"ontgen radius. } \\end{figure}  As we see, interior to a radius $R_{\\rm X}$ the ionization rate due to X-rays exceeds the background level provided by cosmic rays, here assumed at the constant value $\\zeta_{\\rm CR}=10^{-17}$ s$^{-1}$. We call this X-ray dominated region the {\\em R\\\"ontgen sphere}, and $R_{\\rm X}$ the {\\em R\\\"ontgen radius}. In the case of Fig. 1, the average ionization rate inside the sphere of constant density n$_{\\rm H}=10^4$ cm$^{-3}$ is about 20 times that due to cosmic rays. We note that the uncertainty on the empirical value of $\\zeta_{CR}$ is a factor of about five around our assumed value (e.g. van der Tak \\& van Dishoeck 2000).  To gauge the dependence of $R_{\\rm X}$ and $\\zeta_{\\rm X}$ on the model parameters, we have considered X-ray emission spectra at different plasma temperatures, as well as the effects of an absorbing column density of N$_{\\rm H}=5\\times10^{20}$ cm$ ^{-2}$ in front of the X-ray emitting region.  As shown in Table~1, the  R\\\"ontgen radius is insensitive to the input values, while the average ionization rate can vary by a factor of $\\sim$3 (the label $w$ refers to the model with local absorption).  \\begin{table}[h] \\begin{center} \\caption {R\\\"ontgen radius and average ionization rate} \\begin{tabular}{l c| c c| c c}  \\tableline Plasma & kT & $R_{{\\rm X}w}$ & $\\zeta_{{\\rm X}w}$ &  $R_{\\rm X}$ &   $\\zeta_{\\rm X}$  \\\\ model &(keV)&(10$^{17}$ cm)&(10$^{-17}$ s$^{-1}$)&(10$^{17}$ cm)& (10$^{17}$ s$^{-1}$) \\\\ \\tableline  2T\\ & 0.2\\  +\\  1  & 1.5 & 6.3 & 1.6 & 21.0 \\\\ 1T\\ & 1            & 1.4 & 5.1 & 1.5 & 10.5\\\\ 1T\\ & 3            & 1.1 & 4.3 & 1.2 & 7.8\\\\ \\tableline \\tableline \\end{tabular} \\end{center} \\end{table}   \\section {Star formation within a fractal cloud}  The results of the previous section for an individual star can now be extended to an ensemble of young stars formed in a cluster/association. The assumptions that enter in the model are:  \\begin{itemize}  \\item The dense gas distribution has a fractal structure of dimension 2.3 (Falgarone, Phillips \\& Walker 1991). A fractal cloud model is obtained  from hierarchically clustered points. Each point is divided into N=5 random points, all within a distance $1/L^h$ for hierarchical level $h$ and geometric factor $L=2$ (Elmegreen 1997). The fractal dimension is defined as $D=\\log N/ \\log L \\sim 2.3$. In Fig.~2 we display a rendition of the fractal gas distribution, integrated along the line of sight. Here, the gas density varies from 500 cm$^{-3}$ to 10$^6$ cm$^{-3}$ and the smallest structure is 2$^{-6}$ times the characteristic size of the cloud ($\\sim$1.5 pc).  \\begin{figure}[h] \\plotfiddle{LORENZANIA2.eps}{6cm}{0}{40}{40}{-110}{-60} \\caption{Bidimensional representation of the fractal cloud with D=2.3.} \\end{figure}  \\item We assume that at $t=0$ the cloud contains no stars. Star formation then begins and continues for 10 million years. Initially, the rate $\\Psi(t)$ is low and then accelerates exponentially as $\\Psi(t) \\propto e^{\\Delta t/t_{\\rm c}}$, with $t_{\\rm c}=2$~Myr (Palla \\& Stahler 2000).  \\item At any time, star are formed randomly in mass in the interval 0.1--10 M$_\\odot$.  However, we require that after 10 million years their distribution $\\phi(m)$ follows a standard IMF $\\phi(m)\\sim m^{-\\gamma}$, with $\\gamma$ given by, e.g., Scalo (1998).  The behavior of both the star formation rate (SFR) and the mass distribution with time is displayed in the left panel of Figure~3.  \\begin{figure}[h] \\plottwo{LORENZANIA3.eps}{LORENZANIA4.eps} \\caption{ {\\it Left panel}: Time evolution of the SFR and mass distribution of 500 stars born in 10 Myr. {\\it Right panel}: The evolution of the R\\\"ontgen radius in an ambient medium of 10$^4$ cm$^{-3}$ during the PMS contraction of two stars of mass $0.1 {\\rm M}_{\\sun}$ and $1 {\\rm M}_{\\sun}$, respectively. } \\end{figure}  \\item  We assume that the stellar X-ray luminosity is proportional to its bolometric luminosity:  $L_{\\rm X}(M_{\\star},t)=3\\times 10^{-4} {\\rm L}_{\\rm bol}({\\rm M}_{\\star},t)$.  Each star has its associated R\\\"ontgen sphere, computed as in Sect. 2.  Because of the variation of ${\\rm L}_{\\rm bol}({\\rm M}_{\\star},t)$ during PMS contraction, the R\\\"ontgen radius also varies with time, as illustrated in Fig.~3 (right panel).  \\end{itemize} ",
        "conclusions": "The preliminary results described above indicate that the X-rays produced by young stars in a cluster forming cloud may contribute significantly to its ionization degree, especially if the dense gas is not distributed uniformly. Typical R\\\"ontgen radii vary between $\\sim 0.01-0.05$ pc, with an enhancement of the ionization rate of a factor $ \\simeq\\ 5-20$ over the value due to cosmic-rays alone. The mass fraction of a molecular cloud ionized by X-rays and cosmic rays is a function of the total number of stars and of their X-ray luminosity. For a fixed X-ray luminosity of $L_{\\rm X} \\sim 3\\times 10^{-4} L_{\\star}$, the stellar density needed to fill a cloud of 1 pc size is about $\\rho_\\star \\geq 500$~pc$^{-3}$. The required number of stars,  the total X-ray luminosity and the formation efficiency are thus comparable to those observed in cluster forming regions, such as $\\rho$ Ophiuchi. The ionization due to YSOs may therefore increase the coupling of the magnetic field to the cloud and the characteristic time of cloud core collapse via ambipolar diffusion. If   star formation  is indeed {\\em accelerating}, the X-ray emission from YSOs can constitute a natural feedback mechanism to {\\em decelerate} the process. More realistic models (including the effects of the dynamical evolution of the gas during star formation, a distribution of X-ray luminosity and emission spectrum as a function of mass and age, etc.) are being carried out to verify the  initial results."
    },
    "0011/astro-ph0011353_arXiv.txt": {
        "abstract": "We have measured the emission line ratios in a sample of 34 CFRS star-forming galaxies with redshifts between $0.5 < z < 1.0$, and computed their metallicities by means of the empirically-calibrated $R_{23}$ metallicity estimator introduced by Pagel et al.\\ (1979). The current analysis concentrates on the 15 galaxies with $L_{H\\beta} > 1.2 \\times 10^{41}$ erg s$^{-1}$.  Although our results can only be regarded as preliminary until near-IR spectroscopy of H$\\alpha$ and [NII]6583 are available, the metallicities of these galaxies appear to be remarkably similar to those of local galaxies selected in the same way, and there appears to have been little change in the relationship between metallicity and line- and continuum-luminosity from $z \\sim 1$ to today.  At this stage our results do not support the idea that these galaxies, known to be generally small and with late-type morphologies, are dwarf galaxies brightened by large bursts of star-formation, as had been suggested from previous studies. Rather, our findings are more consistent with a picture in which these systems are the progenitors of today's massive metal-rich galaxies.   \\noindent ",
        "introduction": "There is now clear evidence from the Canada-France Redshift Survey (CFRS; Lilly et al.\\ 1995a), and other similar surveys, for evolutionary changes in the field galaxy population over the redshift interval $0 < z < 1$ manifested as an increased number of blue galaxies, with rest-frame $(U-V)_{AB} \\lta 1.4$, and moderate luminosities, $L \\sim L^*$.  The redshift regime $0.5 < z < 1.0$ appears to be an important epoch in the history of the galaxy population.  Evolutionary effects are clearly seen in the galaxy population, and the apparent change of behaviour in the ultraviolet luminosity density of the Universe as a whole around $z \\sim 1$ (Lilly et al.\\ 1996; Madau et al.\\ 1996; Steidel et al.\\ 1999) may indicate that this epoch represents a transition between the ``high redshift'' Universe at $z > 1$ and that seen today.  While the large galaxies appear to be absent at $z \\sim 2$, large spiral galaxies and early-type galaxies are present at $z \\sim 0.8$ in numbers comparable to those seen locally, co-existing with increased numbers of comparably bright but smaller irregular galaxies (Brinchmann et al.\\ 1998, hereafter B98; Lilly et al.\\ 1998; hereafter L98). There is morphological evidence for substantially elevated levels of merger activity at this epoch (Le Fevre et al.\\ 2000).  Much of the work on the $z\\sim1$ field galaxy population has been to date based on gross statistical measures such as the bivariate color-luminosity function (Lilly et al.\\ 1995c; Heyl et al. 1997; Lin et al.\\ 1999).  The HST-based studies of the morphologies and sizes of field galaxies in this redshift regime (e.g. B98; L98), and of the internal kinematics of various subsets of these galaxies (Vogt et al.\\ 1997; Guzman et al.\\ 1997; Mallen-Ornelas et al.\\ 1999) have suggested that the most massive galaxies are evolving relatively slowly (Vogt et al.\\ 1997; L98) while the ``excess'' galaxies in the luminosity function (with $L \\sim L^*$ and blue colors) have generally the irregular morphologies (B98), small sizes (3-5 $h_{50}^{-1}$ Mpc, L98) and low velocity dispersions $\\sigma < 100$ kms$^{-1}$ (Guzman et al.\\ 1997; Mallen-Ornelas et al.\\ 1999) that are usually associated in the local Universe with galaxies 2-3 magnitudes further down the luminosity function. This suggests at first sight quite a strong luminosity brightening in late-type low mass galaxies.  There are still, however, major gaps in our knowledge of galaxies in the crucial $0.5 < z < 1.0$ redshift regime which make it possible that this interpretation of the morphologies and the kinematic data is incorrect.  Not least, the possibility that the small-blue-irregular galaxies are the cores of more massive galaxies ---i.e., the ``down-sizing'' scenario of Cowie et al.\\ (1996)--- cannot be ruled out at this stage. For example, the $K$-band luminosities of typical blue galaxies at $z \\sim 0.8$ are comparable to those of today's massive galaxies, leading to suggestions that at least some of these blue objects may be indeed the progenitors of massive systems.  There has hitherto been almost no systematic study in the $0.5<z<1$ regime of the physical diagnostics that are familiar from studies of the local Universe. For local galaxies, combinations of strong emission lines are routinely used to determine or constrain the nature of the ionizing radiation (Veilleux \\& Osterbrock 1987), the amount of reddening, and the metallicity of the interstellar medium (ISM; Pagel et al.\\ 1979; Kennicutt 1998; Stiavelli 1998; Kobulnicky et al.\\ 1999).  Although a rigorous determination of the ISM metallicity requires knowledge of electron temperature derivable only from intrinsically weak lines, diagnostic line ratios based on stronger lines have been empirically calibrated.  In particular, the $R_{23}$ parameter, $R_{23} = ([OII]3727 + [OIII]4959+5007)/H\\beta$ (Pagel et al.\\ 1979), has been empirically calibrated against metallicity, with an intrinsic scatter of only 0.2 dex. It is a weak function of the ionization ratio [OIII]4959+5007/[OII]3727.  A reversal in $R_{23}$ occurs at $Z \\sim 0.3 Z_{\\odot}$ due to cooling effects, so a low- and a high-metallicity solution are associated with most values of $R_{23}$.  This degeneracy can however be broken using the [OIII]5007/[NII]6584 ratio (Kobulnicky et al.\\ 1999).  Based on data on relatively low redshift galaxies, Kobulnicky et al.\\ have discussed in detail the potential accuracy that could be obtained in using $R_{23}$ to measure metallicities of unresolved galaxies at high redshifts.  These authors examined theoretically the effects of dust reddening, of $H\\beta$ absorption, of spatial averaging over extended galaxies with abundance gradients, and the possible effects of diffuse interstellar gas, and provided prescriptions for dealing with these effects.  Determining the metallicity of the ISM of distant starforming field galaxies is of particular importance, both as a general indicator of the evolutionary state of these systems, and as a constraint on their possible present-day descendants. Some work has been done to determine the metal content of galaxies at redshifts of about $0.1 < z< 0.5$ (Kobulnicky \\& Zaritsky 1999); no information has however been available about the ISM metallicities of galaxies in the $0.5 < z < 1.0$ redshift interval.  We have therefore begun a program of systematic emission line spectroscopy of CFRS galaxies to determine the ISM metallicity of $0.5 < z < 1.0$ field galaxies, and to study how the metal content in these systems correlates with galaxy luminosity, star-formation rate, structure and morphology.  Our program uses the $R_{23}$ metallicity estimator and therefore requires spectroscopy of the [OII]3727, [OIII]4959,5007 and H$\\beta$ lines, which in this redshift range are shifted into the 5000 $\\AA$ - 1 $\\mu$m wavelength region.  Supplementary spectroscopy of H$\\alpha$, [NII]6584 and [SII]6717,6731 lines, which are shifted into the near-infrared J-band, allows determination of the reddening, isolation of active galactic nuclei, and the breaking of the $R_{23}$-degeneracy with metallicity.  In this Letter we report the first results of our program which are based on deep multi-object spectrophotometry over the $0.5 < \\lambda < 1.0 \\mu$m range.  Infrared spectrophotometry has already been acquired for one object, and will be obtained in due course for the remaining galaxies, but the optical data on its own already provides a set of homogeneous data that can immediately be compared at a phenomenological level with the equivalent local sample (see e.g. Jansen et al.\\ 2000). Throughout this paper we use $H_o=50$ Mpc/Km/s and $q_o=0.5$, and refer to solar metallicity $Z_\\odot$ as $12 + Log(O/H) = 8.9$. ",
        "conclusions": "Figures 1a and 1b show the [OIII]5007/[OII]3727 versus $R_{23}$ relation for the local field sample of Jansen et al.\\ (2000) and the high-$L_{H\\beta}$ CFRS galaxies of our sample, respectively.  The galaxies of the Jansen et al.\\ sample were selected from the CfA redshift catalog (Huchra et al.\\ 1983) to span a large range in absolute B magnitude (from -14 to -22), while sampling fairly the changing mix of morphological types as a function of luminosity.  In the figures, the solid and dashed lines are lines of constant metallicity.  From left to right, the dashed lines indicate increasing metallicities from 0.02$Z_\\odot$ to 0.2$Z_\\odot$. At $Z \\sim 0.3Z_\\odot$, the reversal of $R_{23}$ occurs, and the curves of constant metallicity then follow the solid-line sequence in which $Z$ rises up to about 3$Z_\\odot$ from right back to left. The $Z = Z_\\odot$ curve is highlighted with a thicker linewidth.  In order to appropriately compare the local and the high--$z$ samples, the local Jansen et al.\\ galaxies which have L$_{H\\beta} > 1.2\\times10^{41}$ ergs s$^{-1}$ are identified with large circles; local galaxies with smaller $H\\beta$ luminosities are represented by small circles. Furthermore, in Fig 1a, filled symbols represent objects which have a flux ratio $[OIII]5007/[NII]6584 > 2$ and thus metallicities $Z \\lta 0.5 Z_\\odot$, and empty symbols indicate objects with higher metallicities, as indicated by a flux ratio $[OIII]5007/[NII]6584 < 2$ (see Edmunds \\& Pagel 1984). In Figure 1b, the high--$z$ galaxies are represented by the filled squares, with the exception of the two probable-AGN (identified by the asterisks), and of the three objects which may or may not be in the sample on account of their H$\\beta$ upper limits (empty squares).  The fiducial dotted-line box in both panels encompasses the bulk of the Jansen's galaxies.  The arrows on the right-side of the figures represent the direction in which the points of the diagram would shift due to reddening by dust (as described by Cardelli et al.\\ 1989); the length of the arrows refer to an $E(B-V)=0.3$ magnitudes at $z=0.7$. The effects of reddening are a function of the [OIII]5007/[OII]3727 ratio.  Several things are apparent in Figures 1a and 1b.  First, the high redshift galaxies fall in the $R_{23}$ versus [OIII]/[OII] plane in locations that are occupied by galaxies in the local Jansen et al.\\ (2000) sample.  Furthermore, once the AGNs are  excluded, the high redshift sample selected to have $L_{H\\beta} > 1.2 \\times 10^{41}$ erg s$^{-1}$ occupies the same restricted location in the $R_{23}$ versus [OIII]/[OII] plane  as do the objects in the local Universe with similarly high $H\\beta$ luminosities.  At both epochs, galaxies selected to have the same H$\\beta$ luminosities exhibit the same range of $R_{23}$ and [OIII]/[OII]. (We note that there may possibly be a small displacement of the high redshift galaxies towards the upper edge of the dotted box in Figure 1a. This could conceivably be due to the effects of higher dust extinction at high redshifts; it would be premature to claim this  at this stage.)  It is also apparent on Figure 1b that once the two probable-AGNs are again excluded, there are no high-$L_{H\\beta}$ objects at high redshifts (at least in this still small sample) that have $R_{23} \\sim 7$, the value that is non-degeneratively associated with intermediate metallicities, i.e. $Z \\sim 0.3 Z_{\\odot}$. The $Z$-degenerate $R_{23}$ values that are measured for the high-$z$, high-$L_{H\\beta}$ galaxies indicate either rather low ($\\lta 0.1Z_\\odot$) or rather high ($\\sim Z_\\odot$) metallicities for these systems. This is again similar to what is observed in the local sample: high-$L_{H\\beta}$ galaxies avoid the intermediate metallicity regime both at the present and at the $z\\sim 1$ epoch.  The similarity between the high and low redshift samples is further illustrated in Figure 2, which shows the relationship between $R_{23}$ and continuum luminosity M$_B$. There is a rather startling similarity between the high-$z$ and the local galaxies on this diagram.  There is little evidence for an evolutionary change in the relationship between metallicity (as estimated from $R_{23}$) and the line and continuum luminosities for high-$L_{H\\beta}$ field galaxies between $z \\sim 0$ and $z \\sim 0.8$. This extends to higher redshifts the findings of Kobulnicky \\& Zaritsky (1999) which was based on a similar sample at $0.1 < z < 0.5$.  For one of the non-AGN galaxies, CFRS-14.0393, we have already obtained a Keck spectrum in the J-band, and analysed it to derive the intensity of the [NII]6584 emission line.  We will report the details on this analysis elsewhere. The important fact, relevant for the current discussion, is that the [OIII]5007/[NII]6584 ratio that this J-band spectrum has allowed us to measure for this ($z=0.6035$) object undoubtedly places it on the high-metallicity branch of the $R_{23}$ parameter, with a metallicity quite close to the solar value. This one case argues in favour of a high, i.e. about solar metallicity for at least some objects in our $0.5<z<1$ high-$L_{H\\beta}$ sample. The HST morphology of this object is that of a regular two-armed spiral (Schade et al.\\ 1995), and so it is possibly not surprising that this galaxy has a high metallicity. It may be that our high-$L_{H\\beta}$ selection favours large well-formed galaxies relative to the general blue CFRS population. However, there is no indication that this is the case from the HST morphology of three additional objects for which the HST data are available (CFRS-10.1213, CFRS-14.0972 and CFRS-14.1258).  Of course, until we obtain the infrared spectroscopy for the entire sample, the $R_{23}$ degeneracy with $Z$ does not allow us to prove on what branch ---i.e., the low- or the high-metallicity one--- each individual high-$z$ galaxy lies, and the possibility remains that some objects have indeed metallicities $Z\\lta 0.1Z_\\odot$.  However, at this stage, the absence of any galaxies with $Z \\sim 0.3 Z_{\\odot}$ makes this possibility rather contrived, since it would imply a bimodal distribution of metallicities with a ``gap'' around such intermediate values of $Z$.  Therefore, although confirmation must await the infrared spectroscopy, the best working hypothesis seems the one where {\\it all} of the analysed high-$L_{H\\beta}$ high redshift galaxies have relatively high metallicities, i.e. within 40\\% of solar.  The consequences of these findings are interesting in the context of the studies of sizes, morphologies and kinematics of high-redshift galaxies discussed above. It is clear in fact that, at this stage, the metallicity measurements of $0.5<z<1$ systems do not support the idea that many of the small and irregular blue L* galaxies that are responsible for the evolution of the luminosity function back to $z \\sim 0.8$, and which dominate the CFRS at these redshifts (B98; L98) are low mass (i.e. low metallicity) dwarfs brightened by substantial luminosity evolution.  In contrast, the metallicity data seem to suggest the interesting possibility that these small irregular galaxies are in fact the progenitors of today's massive, metal-rich galaxies, but seen in an earlier phase of their evolution when they were already significantly metal-rich but morphologically more disturbed and smaller. Although about a 1/3 of the CFRS has been imaged by HST (see Schade et al 1996; B98), the overlap with this spectroscopic sample is as yet still small.  Testing this idea by studying the morphologies of these metal-rich systems will have important consequences for our understanding of the evolutionary path of massive galaxies.  \\bigskip"
    },
    "0011/astro-ph0011509_arXiv.txt": {
        "abstract": "We investigate the Ly$\\alpha$ line transfer in nearby and high redshift starbursting galaxies, where the effects of high optical depths and the role of dust in the scattering medium are expected to be conspicuous and should be treated in a very careful manner. We present our first results in a dustless, static, and uniform \\ion{H}{1} media with moderate Ly$\\alpha$ line center optical depths $\\tau_0=10^{3-6}$. We assume the temperatures of the media to be $T=10^{1-4}{\\rm\\ K}$, and use a Monte Carlo technique. We investigate the basic processes of the line transfer and confirm the criterion of $a\\tau_0>10^3$ for the validity of the diffusion approximation suggested by Neufeld in 1990, where $a$ is the Voigt parameter. Adopting the model for evolution of a galactic supershell suggested by Tenorio-Tagle et al., we performed calculations on the Ly$\\alpha$ line formation for each evolutionary stage of an expanding supershell. The emergent Ly$\\alpha$ profiles are characterized by the double peaks and the absorption trough at the line center. We found that the absorption troughs expected in most of the evolutionary stages are not wide enough to be observed with current instruments. However, the absorption trough in the Ly$\\alpha$ emission profile from an expanding recombining supershell can be marginally detected. ",
        "introduction": "With the advent of large ground-based and space telescopes, primeval galaxies at high redshifts ($3<z<5$) have been detected through a variety of methods. One of the most efficient detection methods is the Lyman break method which can be used to select high redshifted objects. Subsequent spectroscopy revealed the existence of Ly$\\alpha$ features in either emission or absorption \\citep{ste96,ste99} in many of galaxies showing Lyman break.  As well as being a redshift indicator, Ly$\\alpha$ bears information on the physical state of those galaxies. The emergent Ly$\\alpha$ profile is very sensitive to  various factors including the kinematics, scattering optical depth, and dust content. Hence, it is important to investigate the Ly$\\alpha$ line formation for its cosmological applications.  The Ly$\\alpha$ profiles of many Lyman break galaxies in the early Universe may be classified into three types : (1) symmetric profile, (2) asymmetric or P Cyg type profile, (3) broad absorption profile extending to the Lorentzian wings. It is known that the Ly$\\alpha$ profiles of nearby starbursting galaxies also fall into one of the above three types \\citep{kun98}. It seems to be a general consensus that the P Cyg type Ly$\\alpha$ profile is attributed to the expanding neutral medium surrounding a Ly$\\alpha$ source \\citep{lee98,kun98,ahn00}. This is supported by the existence of the interstellar absorption from low ionization species such as \\ion{O}{1} $\\lambda$1302 and \\ion{C}{2} $\\lambda$1335 which are blueshifted relative to the Ly$\\alpha$ line center \\citep{kun98,spi99}. It is also probable that the profile is formed in galactic superwinds. Assuming that a galactic superwind possesses a Hubble type outflow, \\citet{lee98} calculated the emergent Ly$\\alpha$ profiles and investigated their spectropolarimetric properties.  The P Cyg type Ly$\\alpha$ profile can also originate from a supershell which is a remnant of strong explosions in a central star-forming cluster. From these pieces of observational evidence, \\citet{ten99} proposed a model for the evolution of galactic supershells based upon hydrodynamical simulations and ionization calculations. However, they presented only a schematic description of the Ly$\\alpha$ profiles in each evolutionary stage, so more systematic and quantitative investigations are desired.  Ly$\\alpha$ is one of the strongest resonance lines found in various emission line objects and therefore has been the subject of intensive research for a long time. Both numerical and analytical studies have been carried out by a number of authors. Recent advances in computer technology have made it possible to explore the parameter space that was not previously accessible.  The main objective of this study is therefore to understand the Ly$\\alpha$ line formation in a thick medium by using a sophisticated Monte Carlo method, and to apply the results to the models for the evolution of superbubbles in starbursting galaxies. This is the first paper in a series in which the Ly$\\alpha$ line transfer in an optically thick medium is investigated. In this paper, we consider only moderately thick and dustless media, and defer the cases of extremely thick media with and without dust to subsequent papers.  This paper is composed as follows. In Section 2, we describe the model for starbursting galaxies and our Monte Carlo code. In Section 3, we show our results, and in the following section we discuss the astronomical applications of our works. The final section summarizes our major findings. ",
        "conclusions": "In order to predict the Ly$\\alpha$ profiles from starbursting galaxies, we investigated the Ly$\\alpha$ line transfer in a moderately thick and dustless medium using a Monte Carlo code. Our code is faithful to the quantum mechanics associated with the resonant and wing scattering of Ly$\\alpha$, and the partial frequency redistribution is accurately treated.  We compared our results with the analytic solutions obtained by previous investigators, and found that the results are in excellent agreement. We confirmed the line transfer mechanism such as a single longest flight, a wandering, and a single longest excursion as proposed by \\citet{ada72}. Comparing our emergent profiles with the analytical solution of \\citet{neu90}, we found that the wing approximation does not hold for moderately thick media $10^3<\\tilde{\\tau_0}<10^3/a$, and that it is valid only for the regime $a\\tilde{\\tau_0}>10^3$.  Next, we applied our results to the Ly$\\alpha$ line formation in starbursting galaxies, in relation with the model given by \\citet{ten99}. The following is the summary of the emergent Ly$\\alpha$ profiles for each evolutionary stage of a superbubble in starbursting galaxies. A brief summary is also presented in Table~1, and the spectral capabilities of typical instruments are shown in Table~2 for comparison.  (a) In the early phase of a starbursting galaxy, the galaxy has a static and cold ($T=10\\ {\\rm K}$) envelope, and the line width of the source can be set to be $10\\kms$. We obtained the emergent profiles consisting of double emission peaks and an absorption trough with width $\\Delta\\lambda\\approx 0.01{\\rm\\ \\AA}$ in the rest frame. If this object is redshifted to $z=3$, the width is widened to $\\Delta\\lambda\\approx 0.04{\\rm\\ \\AA}$ in the optical band, which is beyond the observational limit of current instruments.  (b) The surrounding medium subsequently forms a conical \\ion{H}{2} region, and a partially ionized envelope may be formed at the center. We assume that the scattering medium is hot ($T=10^4\\ {\\rm K}$) and the source is at the center of the slab. The emergent Ly$\\alpha$ profiles show symmetric double peaks and an U-shaped absorption trough at the line center.  The trough has a width of $\\Delta\\lambda\\approx 0.2{\\rm\\ \\AA}$, and the features that are formed in nearby starbursting galaxies are found to be observable using the current high resolution far UV spectroscopy such as HST/STIS. When those objects are located at $z=3$, the width of the absorption feature is about $1{\\rm\\ \\AA}$ which can be marginally detected with the Keck LRIS.  (c) In the next stage, an expanding superbubble develops. In this stage, there appear two Ly$\\alpha$ emission peaks. The major peak is formed by the central star cluster and the minor peak is formed by the radiative and expanding supershell. Hence, the major peak is located at the line center, and the minor peak is in the blue part of the line center. According to our investigation, the major and minor peaks may have an absorption trough at the center.  The width of the trough corresponds to $\\Delta \\lambda \\approx 0.25{\\rm\\ \\AA}$. Hence, this central U-shaped absorption feature in nearby starbursting galaxies may be marginally observable using the HST/STIS. When those objects are located at $z=3$, the width of the absorption feature becomes $\\sim 1{\\rm\\ \\AA}$, and this trough should be marginally detected by using ground based large telescopes.  (d) When the neutralization of the supershell is complete, its \\ion{H}{1} column density reaches $N_{HI}=10^{19-22}{\\rm\\ cm^{-2}}$, and its expansion gradually ceases. Hence, we modeled this stage by a hot source surrounded by a static and cold ($T=10\\ {\\rm K}$) scattering medium. Since at this stage the central Ly$\\alpha$ source is extinguished, we considered the ultraviolet continuum source near Ly$\\alpha$. The profiles emergent at this stage are also characterized by double peaks and U-shaped absorption trough.  In this paper, we concentrated on the cases of moderate optical depths. However, the line center optical depth $\\tau_0$ of the neutral envelopes in starbursting galaxies is expected to be $10^{6-9}$. In order for our model to be realistic, we must investigate these extremely thick optical depths and include dust effects, which will enhance the width of absorption trough, and make its detection more feasible. We will study these cases in the future.  In section 3.3 we show the angular distribution $P(\\mu)$ of escaping photons. In the case of a slightly thick medium, we get an almost isotropic distribution given by $P(\\mu)\\propto \\mu$. In a moderately thick medium, photons tend to escape more in the direction parallel to the slab and hence $P(\\mu)$ becomes convex upward. In the extremely thick medium the escaping photons are emergent preferentially in the direction normal to the plane, and $P(\\mu)$ becomes convex downward. This means that when wing scatterings dominate the line transfer, the photons are affected by the geometrical shape of the scattering region and the quantum mechanical properties of scattering. Therefore we expect the development of polarization when Ly$\\alpha$ photons are radiatively transferred in an extremely thick aspherical medium \\citep{cha60,lee98}. This will be considered in the near future."
    },
    "0011/astro-ph0011023_arXiv.txt": {
        "abstract": "Core collapse supernovae (SN) are the final stages of stellar evolution in massive stars during which the central region collapses, forms a neutron star (NS), and the outer layers are ejected. Recent explosion scenarios assumed that the ejection is due to energy deposition by neutrinos into the envelope but detailed models do not produce powerful explosions. There is mounting evidence for an asphericity in the SN which is difficult to explain within this picture. This evidence includes the observed high polarization, pulsar kicks, high velocity iron-group and intermediate-mass elements material observed in remnants, etc.  The discovery of highly magnetars revived the idea that the basic mechanism for the ejection of the envelope is related to a highly focused MHD-jet formed at the NS. Our 3-D hydro simulations of the jet propagation through the star confirmed that the mechanism can explain the asphericities.  In this paper, detailed 3-D models for jet induced explosions of \"classical\" core collapse supernovae are presented. We demonstrate the influence of the jet properties and of the underlaying progenitor structure on the final density and chemical structure. Finally, we discuss the observational consequences, predictions and tests of this scenario. ",
        "introduction": "Supernovae (SN) are among the most spectacular events because they reach the same brightness as an entire galaxy. This makes them good candidates to determine extragalactic distances and to measure the basic cosmological parameters. Moreover, they are thought to be the major contributors to the chemical enrichment of the interstellar matter with heavy elements. Energy injection by SN into the interstellar medium, triggers star formation and feedback in galaxy formation, and is  regarded as a key for our  understanding of the formation and evolution of galaxies.  Core collapse supernovae are thought to be the final stages of the evolution of massive stars which live only $10^6$ to $2 \\times 10^8$ years. Such supernovae could be the brightest objects in the distant past when stars first began to form. A detailed understanding of core collapse is essential to probe the very early phases of the Universe right after the initial star forming period which occurs at redshifts $z \\geq 3 ... 5$. Understanding the mechanism of core collapse supernovae explosions is a problem that has challenged researchers for decades (Hoyle \\& Fowler 1960). In the general scenario for the explosion, the central region of a massive star collapses and forms a neutron star. Eventually, parts of the potential energy will cause the ejection of the envelope. This general scenario  has been confirmed by a wealth of observations including the direct  detection of neutrinos in SN1987A and neutron stars in young supernovae remnants.  In recent years, there has been a mounting evidence that the explosions of massive stars (core collapse supernovae) are highly aspherical. (1) The spectra of core-collapse supernovae (e.g., SN87A, SN1993J, SN1994I, SN1999em) are significantly polarized indicating asymmetric envelopes (M\\'endez et al. 1988; H\\\"oflich 1991; Jeffrey 1991; Wang et al. 1996; Wang et al. 2000). The degree of polarization  tends to vary inversely with the mass of the hydrogen envelope, being maximum for Type Ib/c events with no hydrogen (Wang et. al. 1999, Wang et al. 2000). For supernovae with a good time and wavelength coverage   the orientation of the polarization vector tends to stay constant both in time and in the wavelength. This suggests that there is a  global symmetry axis in the ejecta (Wang et al. 2000). (2) Observations of SN~1987A showed that radioactive material was brought to the hydrogen rich layers of the ejecta very quickly during the explosion (Lucy 1988; Tueller et al. 1991). (3) The remnant of the Cas~A supernova shows rapidly moving oxygen-rich matter outside the nominal boundary of the remnant and  evidence for two oppositely directed jets of high-velocity material (Fesen \\& Gunderson 1997). (4). Recent X-ray observations with the CHANDRA satellite have shown an unusual distribution of iron and silicon group elements with large scale asymmetries in Cas~A (Huges et al. 2000). (5) After the explosion, neutron stars are observed with high velocities, up to 1000  km/s (Strom et al. 1995).   There is a general agreement that the explosion of a massive star is caused by the collapse of its central parts into a neutron star or, for massive progenitors, into a black hole. The mechanism of the energy deposition into the envelope is still debated.  The process likely involves the bounce and the formation of the prompt shock (e.g. Van Riper 1978, Hillebrandt 1982), radiation of the energy in the form of neutrino (e.g. Bowers \\& Wilson 1982)  and the interaction of the neutrino with the material of the envelope and various types of convective motions ( e.g. Herant et al. 1994, Burrows et al. 1995, M\\\"uller \\& Janka 1997, Janka \\& M\\\"uller 1996), rotation (e.g. LeBlanc \\& Wilson 1970, Saenz \\& Shapiro S.L. 1981, M\\\"onchmeyer et al. 1991) and magnetic fields (e.g. LeBlanc \\& Wilson 1970, Bisnovati-Kogan 1971).  Spherically symmetric explosion models rely on the neutrino deposition mechanism. The results depend critically on the progenitor structure, equation of state, neutrino physics,  and implementation of the neutrino transport. Currently,  results are inconclusive even when using sophisticated Boltzman solvers for the neutrino transport. For example, Mezzacappa et al. (2000)  find an explosion whereas Yamada et al. (1999)  and Rammp \\& Janka (2000) do not. Even if successful, these models cannot explain the observed asymmetries. Within the spherical core-collapse picture, additional mechanisms must be invoked which operate within the envelope itself.  Two such mechanisms have been studied in some detail. One is the Rayleigh-Taylor instability which causes mixing of the layers of different composition when the outgoing shock front passes through (M\\\"uller et al. 1989, Benz \\& Thielemann 1990, Fryxell et al. 1991). This effect can explain mixing of the carbon, oxygen and helium-rich layers required for the SN1987A, but none of the simulations were able to account for the  high velocity of Ni observed in SN1987A (Kifonidis et al. 2000). Rayleigh-Taylor mixing provides a rather small-scale structures and can hardly account for the observed polarization which requires a global asymmetry of the expanding envelope ( H\\\"oflich 1991). Another mechanism involves an explosion inside a rapidly and differentially rotating supernova progenitor (Steinmetz \\& H\\\"oflich 1992). With this mechanism it was possible to account for  the polarization in SN1987A which originated from a blue supergiant. This mechanism may have difficulty accounting for the early polarization in  some Type II supernovae (Wang et al. 2000) whose light curves indicate a red-giant progenitors. A strong differential rotation in red supergiants can hardly be expected due to their convective envelopes (Steinmetz \\& H\\\"oflich, 1991).  Attempts have also been made to include multi-dimensional effects into a model of collapse itself. The collapsing core becomes  unstable due to the gradients of both electron mole fraction $Y_e$ and entropy. The developing convection then affects both the neutrino flux and the energy deposition behind the stalled shock. Numerous studies have demonstrated the presence of this effect. It is still debated whether convection combined with the neutrino transport provides the solution to the supernova problem (Rammp et al. 1998 and references therein). In the current calculations, the size and scale of the convective motions seem to be too small to explain the observed asymmetries. The angular variability of the neutrino flux caused by the convection has been invoked to explain the neutron star kicks (Burrows et al. 1995, Janka \\& M\\\"uller 1994). Calculations give kick velocity up to $\\simeq 100$ km/s whereas NS with velocities of several $100$ km/s are common.  Rotation of the collapsing core may also be important. It tends to facilitate the explosion because  the centrifugal barrier reduces the effective potential for the material moving in the equatorial plane and introduces an axial symmetry in the fluid motions. Simulations made so far indicate that the rotation alone has no or has only a weak effect on the explosion (e.g. M\\\"onchmeyer et al. 1991, Zwerger \\& M\\\"uller 1997). In the latter case it induces a rather weak asymmetry of the explosion with more energy going along the rotational axis.  It has long been suggested that the magnetic field can play an important role in the explosion (LeBlanc \\&  Wilson 1970; Ostriker \\& Gunn 1971, Bisnovati-Kogan 1971, Symbalisty 1984). LeBlanc and Wilson simulations showed the amplification of the magnetic field due to rotation and the formation of two oppositely directed, high-density, supersonic jets of material emanating from the collapsed core. Their simulations assumed a rather high initial magnetic field $\\sim 10^{11}$ Gauss and produced a very strong final fields of the order of $\\sim 10^{15}$ Gauss which seemed to be unreasonable at the time. The recent discovery of pulsars with very high magnetic fields (Kouveliotou et al. 1998, Duncan \\& Thomson 1992) revives the interest in the role of rotating magnetized neutron stars in the explosion mechanism. It is not clear whether a  high initial magnetic field required for the LeBlanc \\&  Wilson mechanism is realistic. On the other hand it may not be needed. Recently Mayer and Wilson (2000, private communication) have suggested that the field amplification and the generation of the jet may continue during the first few seconds of the cooling of the neutron star when the neutron star shrinks rapidly. The current picture of the core collapse process is unsettled. A quantitative model of the core collapse must eventually include all the elements mentioned above.  Due to the difficulty of modeling core collapse from first principles, a very different line of attack on the explosion problem has been used extensively and proved to be successful in understanding of the supernova problem, SN1987A in particular (Arnett et al. 1990, Hillebrand \\& H\\\"oflich 1991). The difference of characteristic time scales of the core (a second or less) and the envelope (hours to days) allows one to divide the explosion problem into two largely independent parts - the core collapse and the ejection of the envelope. By assuming the characteristics of the energy deposition into the envelope during the core collapse, the response of the envelope can be calculated. Thus, one can study the observational consequences of the explosion and deduce characteristics of the core collapse and the progenitor structure. This approach has been extensively applied in the framework of the 1D spherically symmetric formulation. The major factors influencing the outcome have been found to be the explosion energy and the progenitor structure. The same approach can be applied in multi-dimensions to investigate the effects of asymmetric explosions. In this paper we study the effects and observational consequences of an asymmetric, jet-like deposition of energy inside the envelope of a core-collapse supernova.  \\noindent ",
        "conclusions": "We have numerically studied the explosion of Core Collapse supernovae caused by supersonic jets generated in the center of the supernova as  a result of the core collapse into a neutron star. We  simulated  the process of the jet propagation through the star, and the redistribution of elements. A  strong explosion and a high efficiency for the conversion of the jet energy requires low jet velocities or a low, initial collimation of the jet.  With increasing extension of the envelope, the conversion factor increases. Typically, we would expect higher kinetic energies in SNe~II compared to SNe~Ib/c if a significant amount of explosion energy is carried away by  jets. Within the framework of jet-induced SN, the lack of this evidence suggests that the jets  have low velocities.   For the compact progenitors of SNe~Ib/c, the final departures of the iso-density contours from sphericity are typical a factor of two. This will produce a linear polarization of about 2 to 3 \\% (Fig. 6) consistent with the values observed for Type Ib/c supernovae. In case of a red supergiant, i.e. SNe~II, the   asphericity is restricted to the inner few solar masses. In the latter case, the iso-densities show an axis ratios  of up to $\\approx$ 1.4 at the innermost, hydrogen-rich layers.  The outer layers remain spherical. This has strong  consequences for the observations,  in particular, for polarization measurements. In general, the polarization should be larger in SNe~Ib/c compared to classical SNe~II which is consistent with the observations by Wang et al. (2000). Early on, we expect no or little polarization in supernovae with a massive, hydrogen rich envelope which will increase with time to about 1 \\% (H\\\"oflich 1991), depending on the inclination  the supernovae is observed. This is also consistent both with the long-term time evolution of SN1987A (e.g. Jefferies 1991) and, in particular, the plateau supernova 1999em which has been observed recently with VLT and Keck (Wang et al. 2000; Leonard et al. 2000).  The He, C, O and Si rich layers                     of the progenitor show   characteristic, butterfly-shape structures. This  overall morphology and pattern should be observable in supernovae remnants, e.g. with the Chanda observatory despite some modifications and instabilities when the expanding medium interacts with the interstellar material.  During the explosion, the jets bring heavy and intermediate mass elements into the outer layers including $^{56}Ni$. Due to the high entropies of the jet material close to the center, this may be a possible site for r-process elements. Spatial distribution of the jet material will influence the properties of a supernova. In our model for a SN~II, the jet material stalled within the expanding envelope corresponding to a velocity of $\\approx 4500 km/sec$ during the phase of homologous expansion. In SN1987A, a bump in spectral lines of various elements has been interpreted by material excited by a clump of radioactive $^{56}Ni$ (Lucy 1988). Within our framework, this bump may be a measure of region where the jet stalled. This could also explain the early appearance of X-rays in SN1987A which requires strong mixing of radioactive material into the hydrogen-rich layers (see above). We note that, if this interpretation is correct, the 'mystery spot' (Nisenson et al. 1988) would be unrelated. In contrast to 1-D simulations, we find in our models strong, continuous fallback over an extended period of time, and a lack of an inner, almost empty cavity. This significant amount of fallback and the consequences for the secondary formation of a black hole shall be noted. Moreover, fallback and the low velocity material may alter the  escape probablity for $\\gamma $-rays produced by radioactive decay of $^{56}Ni$. In general, the lower escape probability is unimportant for the determination of the total $^{56}Ni$ production by the late LCs   because full thermalization can be assumed in core collapse SN during the first few years. However, in extreme cases such as SN98BW (e.g. Schaefer et al. 1999), only a small fraction of gamma's are trapped. Effects of multi-dimensionallity will strongly alter the  energy input by radioactive material and disallow a reliable estimate for the total $^{56}Ni$ mass.  Finally, we want to emphasize the limits of this study and some of the open questions which will be addressed in future. We have assumed that jets are formed in the course of the formation of a neutron star, and have addressed observational consequences and constrains. However, we have not calculated the jet formation, we do not know if they really form, and, if they form, whether they form in all core-collapse supernovae. We cannot claim that the jets are the only mechanism that can explain asphericity in supernovae although we are not aware of the others. Qualitatively, the observational properties of core collapse supernovae are consistent with jet-induced supernovae and support strongly that the explosion mechanism is highly aspherical  but no detailed comparison with an individual object has been performed.   \\noindent{\\sl Acknowledgments:} We want to thank our collegues for helpful discussions, in particular, E.S. Oran, L. Wang, J.C. Wheeler,  Inzu Yi A., C. Mayers, J.C. Wilson, A. Chieffi, M. Limongi, and O. Straniero.    This work is supported in part by NASA Grant LSTA-98-022."
    },
    "0011/astro-ph0011215_arXiv.txt": {
        "abstract": "Integral field spectroscopy of the central parsec of the Galactic Center was obtained at 2.06~$\\mu$m using BEAR, an imaging Fourier Transform Spectrometer, at a spectral resolution of 74 km~s$^{-1}$. Sixteen stars were confirmed as \\lq\\lq helium stars\\rq\\rq\\/  by detecting the \\ion{He}{I} 2.058~$\\mu$m line in emission, providing a homogeneous set of fully resolved line profiles. These observations allow us to discard some of the earlier detections of such stars in the central cluster and to add three new stars. The sources detected in the BEAR data were compared with adaptive optics images in the K band to  determine whether the emission was due to single stars. Two sub-classes of almost equal number are clearly identified from the width of their line profiles, and from the brightness of their continuum. The first class is characterized by very broad line profiles (FWHM $\\simeq$ 1000 km~s$^{-1}$) and by their relative faintness.  The other, brighter in K by an average factor of $\\sim$ 9, has a much narrower emission component of width $\\simeq$ 200 km~s$^{-1}$.  Most of the emission lines show a P Cygni profile. From these results, we propose that the latter group is formed of stars  in or near the LBV phase, and the other one of stars at the WR stage.  The division into two groups is also shown by their spatial distribution, with the narrow-line stars in a compact central cluster (\\object{IRS~16}) and the  other group distributed at the periphery of the central cluster of hot stars. In the same data cube, streamers of interstellar helium gas are also detected. The helium emission traces the densest parts of the SgrA West Mini-Spiral.  Several helium stars have a radial velocity comparable to the velocity of the interstellar gas in which they are embedded. In the final discussion, all these findings are examined to present a possible scenario for the formation of very massive stars in the exceptional conditions of the vicinity of the central Black Hole. ",
        "introduction": "The very inner region of the Galactic Center (GC) is the focus of many studies as it offers the unique opportunity to study star formation and the extreme gas kinematics peculiar to the vicinity of a 2.5$\\times$10$^{6}$ M$_{\\odot}$ black hole (Genzel et al. \\cite{genzel97}, Ghez et al. \\cite{ghez}). The presence of an unusually broad 2.058~$\\mu$m neutral helium line in emission was among the early known  peculiarities of the central infrared source, originally called \\object{IRS~16}  (Hall et al. \\cite{hall}). Continuously improved spatial resolution has made it possible to tie this emission to individual stars and to suggest that this emission is explained by the presence of massive, young, hot stars (Najarro et al. \\cite{najarro97a}). However, even if the formation of high mass stars  was favored in the GC (Morris \\cite{morris93}), the prediction of evolving starbursts cannot fully explain the large abundance of massive, emission-line stars which are normally very  rare and short-lived (Lutz \\cite{lutz}). Therefore, more spectroscopic observations are warranted to better understand the unique conditions in the central parsec of the Milky Way which can lead to the formation of numerous helium emission-line stars. An exact census and a precise determination of the physical properties of these stars is also important since they should significantly contribute to the ionization of the central parsec.  In this paper, we present new data obtained with an original type of integral field spectrometer, an imaging Fourier Transform Spectrometer called BEAR, on the 3.6-m Canada-France-Hawaii Telescope. The use of this instrument represents an effort to make a significant step in associating the best possible spatial resolution and a high spectral resolution in the near infrared. The spatial resolution is not limited by a slit width, as with a standard grating spectrometer. It corresponds to the common seeing conditions at the CFH Telescope on Mauna Kea at 2~$\\mu$m ($\\simeq$ 0.6$\\arcsec$). The spectral resolution is provided by the FTS. To complete the star detection, an adaptive optics (AO) image of the same field in the K band was utilized. The processing of the BEAR data cube is described in Sect.~\\ref{process}. All the new results obtained from this study are presented in Sect.~\\ref{result}, including the display of the \\ion{He}{I} 2.058~$\\mu$m line profiles of all the detected stars and, for the first time, the mapping of flows of interstellar helium. A detailed review of the detected stars follows in Sect.~\\ref{notes}. Finally, a discussion of the nature of the \\ion{He}{I} stars, of the link between these stars and the helium flows, and of a possible star formation scenario are presented in Sect.~\\ref{discuss}. ",
        "conclusions": "New results on the population of \\ion{He}{I} stars in the inner region of the Galactic Center have been obtained with the firm indication of two classes of massive hot stars, which suggest a formation in a disk of gas around SgrA$^{\\star}$. This analysis is based on a single line, the \\ion{He}{I} 2.058~$\\mu$m line. Of course, while that is not sufficient to establish a complete spectral classification, one can nonetheless consider that this study represents a necessary initial selection test, since it has been shown that several stars in previous studies were wrongly considered as \\ion{He}{I} stars. By consequence, it illustrates also the risk of false detections in the search for emission-line stars toward the inner Galaxy simply by using narrow-band photometry centered on the \\ion{He}{I} 2.058~$\\mu$m line, or other near-infrared emission lines. The importance of high spectral resolution combined with high spatial resolution is paramount for distinguishing stellar and interstellar emission. Similar data obtained with BEAR already exists on Br$\\gamma$ (Morris \\& Maillard \\cite{morris99}). With the same effort for separating the stellar and the interstellar component, the Br$\\gamma$ line profile for the 16 confirmed \\ion{He}{I} stars should be retrieved, making another test on this stellar population. The next goal is a larger spectral coverage to complete the spectral criteria. That has already been  done by several previous works, but always at medium resolution and with slit spectrometers, for which the source confusion is not easily controlled. Only a slitless technique like that employed by BEAR makes this control possible. Also, similar studies on other critical lines must be conducted. However, the next major step, instead of spectro-imaging at seeing-limited resolution as we have presented, and tried to improve by combining with AO imaging, will be infrared spectro-imaging at the diffraction-limited spatial resolution of a large telescope. That is the only way to detect more sources in order to remove all the identification ambiguities. As illustrated here, this should be combined with a spectral resolution of at least 5000, which is not an easy goal.  Finally, it would be important to conduct similar studies in other stellar clusters like the \\lq\\lq Arches\\rq\\rq\\/ and  \\lq\\lq Quintuplet\\rq\\rq\\/ clusters (Figer et al. \\cite{figer99a}, \\cite{figer99b}) where the identification of LBV stars has also been proposed, notably the Pistol star (Figer et al. \\cite{figer99c}), to determine comparatively the conditions of evolution of young compact clusters at Galactocentric distances well beyond the central pc of the Milky Way."
    },
    "0011/astro-ph0011573_arXiv.txt": {
        "abstract": "First results from  high-resolution coronal spectroscopy with the {\\it XMM-Newton} Reflection Grating Spectrometers (RGS) are reviewed. Five stellar systems (HR~1099, Capella, Procyon, YY Gem, AB Dor) have been observed. The emphasis of the present paper is on overall coronal structure. Elemental abundances in {\\it active stars} are found to be `anomalous' in the sense that they tend to increase with increasing First Ionization Potential (FIP - i.e., signifying an inverse FIP effect). Coronal densities are measured at levels of a few times $10^{10}$~cm$^{-3}$ for cooler plasma, although there are indications for very high densities in the hotter plasma components. ",
        "introduction": "\\begin{figure}[t!] \\plotone{mguedel_1.eps} \\caption{Five spectra observed with RGS, smoothed to 30~m\\AA\\ effective resolution. The average coronal temperature increases from Procyon (top), to Capella (2nd panel), YY Gem (3rd panel), AB Dor (4th panel) and HR~1099 (bottom).} \\end{figure}  High-resolution X-ray spectroscopy now available from {\\it Chandra} and {\\it XMM-Newton} has opened new access to coronal diagnostics. For the first time, densities, elemental abundances, the thermal structure, and mass motions can be measured explicitly based on atomic X-ray line diagnostics  for a large number of sources. Such measurements are pivotal for our understanding of coronal structure, coronal heating, and mass transfer from the lower layers into the corona.  The two Reflection Grating Spectrometers (Brinkman et al., this volume) onboard {\\it XMM-Newton} (Jansen et al., this volume) have started their mission with the First-Light target HR 1099 and later continued to obtain deep X-ray spectra of various coronal sources for an extended survey of coronal activity. We report here a first summary of results (see also Audard et al., this volume).  \\begin{figure}[t!] \\plotone{mguedel_2.eps} \\caption{Extract from Fig. 1, showing the Fe L-shell region, together with lines of Ne and O.} \\end{figure} ",
        "conclusions": "Stellar coronae have been compared with scaled-up versions of the solar corona although this analogy has clearly become doubtful. Not only are the X-ray luminosities of active stars orders of magnitude larger than the Sun's, but the thermal structure is  also very different, ranging from cooler solar-like structures of a few MK (probably analogous to the Procyon corona) up to {\\it non-flaring} coronal components at 20$-$30~MK. Spectroscopy clearly provides the necessary information for further study.  The most crucial results from this early {\\it XMM-Newton} RGS stellar coronal survey set active stars apart from the Sun. Coronal elemental abundances, although known before to deviate from solar behavior, are now showing a clear systematic dependence on the FIP value, such that high-FIP elements (O, N, Ne) are largely overabundant compared to low-FIP elements (Fe, Mg, Si, Ni), thus defining an {\\it inverse FIP effect}, contrary to the solar trends. The Sun does occasionally reveal Ne- and S-rich flares (i.e., high-FIP elements are enriched in the flaring plasma), and analogous processes may be responsible for the anomalously high Ne abundance seen in active stars.  Densities found so far for active stars are high for solar standards, even for relatively cool plasma. High densities in hot plasma, although still only marginally detected, have no solar analogy except for flares. Taken together, we speculate that the mechanisms for heating to high temperatures and for inverting the coronal abundance pattern are related, but unknown on the {\\it non-flaring} Sun."
    },
    "0011/astro-ph0011090_arXiv.txt": {
        "abstract": "We discuss how several of the questions that remain unclear on the physics of Active Galactic Nuclei may find elements of answers when using in the coming years the extraordinary set of instruments that will be flying simultaneously to observe in all bands of the electromagnetic spectrum. The choice of questions mentioned here is personal and not exhaustive. ",
        "introduction": "Wondering what may be learned about Active Galactic Nuclei (AGN) in the coming years requires that one thinks about what is currently known and about the type of knowledge that is relevant. Indeed, AGN seem to be complex systems that need  not be described at a level at which the general physical processes are hidden behind local and/or statistical fluctuations. The learning process may in some sense be similar to understanding  the physics of the terrestrial atmosphere by doing meteorological observations. The aim in this case is not to describe the  local weather but to extract from the complex descriptions of the observations the guiding principles. Clearly, what knowledge or understanding one wishes to obtain on any given subject is a matter of personal taste. Indeed many people, even among astronomers, might have only a limited interest in AGN. The following account is therefore a personal approach to the subject.  \\vspace{1cm} Among the facts that are well established now and documented in a number of books and reviews in recent years one finds:  \\begin{description} \\item[The source of energy] is the gravitational energy of matter falling into a massive black hole.  With this general principle one deduces two  characteristics of AGN, namely an estimate of the mass of the central black hole through the Eddington luminosity (more appropriately, this estimate gives a lower boundary to these masses). This mass  is found to be up to $10^{10}$ solar masses. Also from the luminosity of the objects and from reasonable assumptions on the efficiency with which accreted matter can be made available to be radiated one deduces the rate at which matter must be accreted into the central black hole. Rates of up to 100 solar masses per year for the brightest nuclei are found.  The role of the rotation of the central black hole in the global energetics still needs be understood.  \\item[The emission processes] through which energy is radiated are now mostly identified. They include synchrotron emission, emission of dust heated by ultraviolet and X-ray radiation fields, Compton processes (be they thermal or due to relativistic electron distributions) and emission from optically thick media (black body radiation). It is  possible from the identification of the emission processes and the shape of the spectral energy distributions to deduce (some of) the physical conditions of the emitting material. \\item[The state of the gas] is known for the gas that emits the lines charateristics of most AGN types. Temperature, density and filling factors as well as the origin of the ionisation can be read from the spectra and provide a picture in which photoionisation plays a dominant  role. \\item[Aberrations and geometry] play  an important role through Doppler boosting of components related to relativistic jets and through absorption of some components that happen to lie behind dense clouds of dust and gas. These effects considerably modify the appearance of AGN and probably lead to our AGN classification scheme. \\item[AGN sit in the center of galaxies] whith which they must have intense relations, be it only through material that is falling into the galactic potential well to the AGN. \\item[The evolution] of AGN is well documented. The AGN  number density and/or luminosity have greatly changed over the age of the Universe. \\item[AGN have a very rich phenomenology] for which we have large amounts of data. We thus have descriptions of many spectral energy distributions, of many types of variations in different objects and at different wavelengths and we have a vast zoology of objects in which populations differ greatly in one or several characteristic parameters. We are not always succesful in understanding this rich phenomenology in terms of underlying physical processes. \\end{description} ",
        "conclusions": ""
    },
    "0011/astro-ph0011059_arXiv.txt": {
        "abstract": "We used the near infrared spectrometer ISAAC at the ESO {\\it Very Large Telescope} to map the velocity field of Centaurus A (NGC 5128) at several position angles and locations in the central 20\\arcsec\\ of the galaxy. The high spatial resolution ($\\sim 0\\farcs5$) velocity fields from both ionized and molecular gas (\\PB, \\FeII, \\BG\\ and \\HMOL) are not compromised by either excitation effects or obscuration. We identify three distinct kinematical systems: (i) a rotating \"nuclear disk\" of ionized gas, confined to the inner 2\\arcsec, the counterpart of the \\PA\\ feature previously revealed by HST/NICMOS imaging; (ii) a ring-like system with a $\\sim 6\\arcsec$ inner radius detected only in \\HMOL\\, likely the counterpart of the 100\\PC-scale structure detected in CO by other authors; (iii) a normal extended component of gas rotating in the galactic potential. The nuclear disk is in keplerian rotation around a central mass concentration, dark ($M/L> 20\\,\\Mo/\\LKo$) and point-like at the spatial resolution of the data ($R<0\\farcs25\\sim 4\\PC$). We interpret this mass concentration as a supermassive black hole. Its dynamical mass based on the line velocities and disk inclination ($\\I>15\\arcdeg$) is $\\M = 2^{+3.0}_{-1.4}\\xten{8}\\Mo$. The ring-like system is probably characterized by non-circular motions; a \"figure-of-8\" pattern observed in the \\HMOL\\ position-velocity diagram might provide kinematical evidence for the presence of a nuclear bar. ",
        "introduction": "Mass accretion onto a supermassive black hole (BH) is widely accepted as the most likely mechanism for powering Active Galactic Nuclei (AGN).  It provides a simple and efficient way to produce the necessary large amount of energy in a small volume. Demonstrating the existence of supermassive black holes and measuring their masses has thus been a crucial element in confirming the standard view of AGNs.  Over the last few years, there have been direct dynamical mass determinations for BHs in nearby galaxies in sufficient numbers and with sufficient accuracy to enable detailed demographic studies \\citep[e.g. the recent reviews by][]{macchetto99,ho98,ford97,marel97,kr95}. The possible correlation observed between the mass of the BH and that of the host spheroid (either the bulge of a spiral galaxy, or the entire mass of an elliptical galaxy) has prompted serious studies of the relationship between BH growth and host galaxy evolution \\citep{magorrian,rich98}.  The nuclei of many AGNs and normal galaxies alike are heavily obscured by dust, making it difficult to study kinematics at visible wavelengths. For example, the 3CR HST snapshot survey has shown that the nuclei of radio galaxies are often covered by foreground dust lanes \\citep{dekoff96,martel99}.  This has been further confirmed by other sensitive surveys with WFPC2 and NICMOS \\citep{malkan,regan}.  Such dust obscuration can distort both absorption line and emission line velocity field measurements to such an extent that reliable BH mass estimates are difficult or impossible using optical data alone.  IR spectroscopy is thus increasingly important to continue to find and measure the masses of black holes in active galaxies.  Centaurus A (NGC5128), the closest active giant elliptical galaxy, nicely illustrates the problem of dust obscuration. At optical wavelengths, its nucleus is made almost completely invisible by its dust lane, rich in molecular and ionized gas, which obscures the inner half kiloparsec of the galaxy.  A comprehensive recent review on the properties of the galaxy and of its AGN is presented by \\citet{israel}.  The kinematics of the ionized gas was mapped in \\HA\\ by \\citet{bta87} and subsequently modeled by \\citet[hereafter NBT92]{nbt92} who showed that the dust lane can be kinematically and morphologically explained with a thin warped disk geometry.  On the assumption of closed circular orbits, the \\HA\\ velocity field is consistent with the rotation curve expected from the galactic potential.  Kinematical mapping in CO by \\citet[hereafter Q92]{q92} confirmed the previous results, showing that the material observed in CO is dynamically and geometrically very similar to that seen in \\HA.  \\citet{q93} demonstrated that the thin warped disk model can explain the observed near-IR morphology.  \\citet{schreier96} showed that HST R-band imaging polarimetry is also consistent with dichroic polarization from such a disk.  The first observational evidence for a large mass concentration at the nucleus of Centaurus A came from CO kinematical observations of the nuclear region; these suggested a compact, $100-300\\PC$-scale circumnuclear ring enclosing a mass of $\\sim\\ten{9}\\Mo$ \\citep{israel90,israel91,ryd93,hawa93}. ISO imaging by \\citet{mira99} provided morphological evidence that the molecular material within the central 2\\KPC\\ might be in the form of a gaseous bar. \\PA\\ narrow band images obtained with NICMOS on the HST \\citep[hereafter Paper I]{schreier98}, revealed an elongated ($\\sim 2\\arcsec\\times1\\arcsec$) emission line region centered on the nucleus; this was interpreted as an extended accretion disk around the AGN. \\citet[Paper II]{marconi00} presented a detailed study of the photometric structure of the galaxy interior to the dust lane using NICMOS and WFPC2 which led to a black hole mass estimate of $\\ten{9}\\Mo$, using BH growth models by \\citet{marel99}. To confirm the presence of a black hole, and to better measure its mass, requires high spatial and spectral resolution IR spectroscopy, as is now becoming possible with the new generation of 8m-class telescopes, with state of the art infrared arrays.  In this paper, we report the results of a high spatial resolution near-IR spectroscopic study of both the ionized and molecular gas in the central region of the galaxy.  Our data, obtained with ISAAC at the ESO {\\it Very Large Telescope}, allow us to study the kinematics of the circumnuclear disk discovered in Paper I on sub-arcsecond scales and to obtain a dynamical mass estimate for the central black hole.  In Section \\ref{sec:obs} we discuss the spectroscopic observations and data reduction. The results themselves are described in Section \\ref{sec:res}, which also discusses the difference in the kinematics between the ionized and the molecular gas, and compares our results with those obtained previously at other wavelengths.  Section \\ref{sec:fit} presents our kinematic modeling of the ionized gas velocity field of the nuclear disk. In Section \\ref{sec:discuss}, we summarize the evidence for a supermassive BH at the nucleus of Cen A and the consequences for understanding AGNs and the BH-mass vs bulge-mass relationship. In particular, in Section \\ref{sec:gasbar} we discuss the kinematics of the molecular gas and provide direct evidence for the kinematic signature of a weak gaseous bar.  Our conclusions are presented in Section \\ref{sec:summary}.  Throughout the paper, we assume a distance to Cen A of 3.5\\MPC\\ \\citep{tonry91,hui,soria96} and an average heliocentric systemic velocity  of $V_{Hel}=543\\pm 2\\KM$ \\citep{israel}. ",
        "conclusions": "Discussion}  The observations presented above provide strong evidence for a dark mass concentration, very likely a supermassive black hole, at the nucleus of Cen A. The BH sphere of influence has an outer radius \\begin{equation} r_\\bullet\\simeq \\frac{G\\M}{\\sigma^2} = 11\\PC \\left(\\frac{\\M}{\\ten{8}\\Mo}\\right) \\left(\\frac{\\sigma}{200\\KM\\S\\1}\\right)^{-2} \\end{equation} where $G$ is the gravitational constant and $\\sigma$ is the velocity dispersion of the stars.  In Cen A, the velocity dispersion has a mean value of $\\simeq 120\\KM\\S\\1$ between 40\\arcsec\\ and 80\\arcsec\\ from the centre with a nuclear value of $\\sim 150\\KM\\S\\1$ \\citep{wilk86}.  With the caveat that extinction could lead to an underestimate of the nuclear value, the radius of influence of the BH would then be $r_\\bullet\\sim 20 M_8 \\PC$ i.e. $\\sim 1.1\\arcsec M_8$. Indeed the nuclear rotation curves within 2\\arcsec\\ of the nucleus do not show any evidence of an extended mass distribution. The radius of influence of the BH could also be estimated as the radius where the enclosed stellar mass equals \\M.  From Fig. \\ref{fig:nuker}c, for $\\Upsilon=1 \\Mo/\\LKo$ this takes place at $R\\sim 1\\arcsec$, or 17pc, (for \\M=\\ten{8}\\Mo), consistent with the above estimate.  The presence of a supermassive BH at the nucleus of Cen A has implications for the energetics of the active nucleus and for the general BH-mass vs bulge-mass relationship.  \\subsection{BH mass and AGN activity}  In the 2-10\\KEV\\ range, the nucleus of Centaurus A has an unabsorbed X-ray luminosity of 6\\xten{41}\\ERG\\S\\1\\, with a power law spectrum typical of normal AGNs - $\\Gamma=1.96\\pm0.1$ \\citep{turner97}. For a typical quasar, $L_X(2-10\\KEV)/L_\\mathrm{bol}\\sim 0.03$ \\citep{elvis}; assuming that Cen A has a similar Spectral Energy Distribution (SED), its expected bolometric luminosity is $\\simeq 2\\xten{43}\\ERG\\S\\1$, in agreement with other estimates presented in literature \\citep[e.g.][]{kinzer}.  This value indicates that Cen A is radiating well below its Eddington luminosity as $L_\\mathrm{bol}/L_{Edd}\\simeq 2.4\\xten{-3} M_8^{-1}$, where $M_8^{-1}$ is the BH mass in units of \\ten{8}\\Mo.  This is similar to the case of M87, another well-studied radio galaxy characterized by a prominent jet and a supermassive BH \\citep[c.f.][]{m87bh}.  The HST survey of nuclei of FR~I radio-galaxies by \\citet{chiab} indicates that sub-Eddington luminosities are typical of this class of sources to which Cen A belongs.  Luminosities much below the Eddington limit are the result either of a very low accretion rate or of an inefficient radiation process.  In the latter case it has been suggested that the accretion is dominated by advection \\citep[Advection Dominated Accretion Flow - ADAF; e.g.][]{narayan} in which case the continuum emitted by the AGN lacks the UV \"big-blue bump\" emitted by the standard thin accretion disk \\citep[e.g.][]{kup00,dimatteo}.  Using ISO LWS spectra, \\citet{unger00} have shown that the total FIR luminosity of Cen A in a 70\\arcsec\\ beam centered on the nucleus is $\\LFIR=3.2\\xten{9}\\Lo$, comparable to the luminosity expected from the AGN, $\\LAGN\\simeq5\\xten{9}\\Lo$, if it has the typical spectrum of a quasar.  If the latter assumption is true, the sub-Eddington luminosities are the consequence of a low accretion rate.  However, since FIR emission lines have excitation and equivalent widths typical of starburst galaxies, \\citet{unger00} suggest that the total FIR luminosity of Cen A is dominated by star formation instead of AGN activity.  In this case the bolometric correction applied to the AGN X-ray emission is overestimated by a large factor (a factor of 10 would imply $\\LAGN/\\LFIR\\simeq 0.15$) suggesting that the big blue bump is weak or absent as it happens in the ADAF regime.  Indeed, in Paper II we showed that the extrapolation of the X-ray power law to the optical and near IR can fully account for the observed nuclear emission.  \\subsection{The BH-mass vs bulge-mass relationship}  There is increasing evidence that black holes exist in most or all galaxies, and it has been suggested that a rough correlation exists between the mass of the black hole and the mass of the bulge of the host galaxy \\citep[e.g.][]{kr95,rich98}.  In this framework, we can compare the central dark mass concentration in Cen A with the total mass of the galaxy.  The total mass of Cen A has been derived by photometrical and kinematical studies. \\citet{hui} and \\citet{mathieu} find $M_\\mathrm{sph}=(4\\pm1) \\xten{11}\\Mo$. However, in the BH-bulge mass relationships, the spheroid mass is usually determined by assuming a mass-to-light ratio from the spheroid luminosity in a certain band (usually V or B).  To compare Cen A with existing BH-bulge relationships, we thus have to verify that the photometry results are compatible with the kinematics.  Indeed, the total V mag of Cen A is V=5.75mag (from \\citet{dufour}, after correcting for the nuclear light profile - Paper II), which corresponds to \\MV=-21.95mag. Therefore $\\log (L/\\LVo) = 0.4(-\\MV+4.83)=10.7$.  For normal galaxies, $M/\\LV=5(L/\\ten{10}\\LVo)^{0.2}$ in solar units, and therefore the expected mass of Cen A is $M\\sim 3\\xten{11}\\Mo$ in good agreement with the kinematical determination.  For Cen A, then: \\begin{equation} \\frac{\\M}{M_\\mathrm{sph}}\\simeq 2.5\\xten{-4} M_8 \\end{equation} where $M_8$ is the BH mass in units of \\ten{8}\\Mo.  The well known BH-mass vs bulge-mass relationship \\citep{rich98} shows that $\\log(\\M/M_\\mathrm{sph})$ has a roughly gaussian distribution, with an average of -2.27 (i.e. a 0.005 ratio) and standard deviation of 0.5.  The expected relationship from Quasar light is $\\simeq 6\\xten{-4}$, a factor ten lower.  This apparent discrepancy suggests that AGN accretion is less efficient than previously thought or that most of AGN activity is obscured \\citep{fabian}.  Note however that the large observed value of 0.005 is based on the \\citet{magorrian} paper where, given the simplified assumptions on the stellar systems and the low spatial resolution of the spectroscopic data, BH masses are likely overestimated. Indeed Cen A is almost 2.5$\\sigma$ (in log) below the average of the Richstone et al.~and Magorrian et al.~papers while it is in good agreement with the expectation from Quasar light. The observed $\\M/M_\\mathrm{sph}$ for Cen A is also in agreement with the same ratio determined by \\citet{wandel} for Seyfert 1 galaxies from reverberation mapping.   \\subsection{\\label{sec:gasbar}The \\HMOL\\ \"ring\"}  We have showed that while the \\HMOL\\ rotation curve well fits our picture of a fast rotating disk around a SMBH within the central 2\\arcsec, at larger radii the \\HMOL\\ velocities at NUC-2 are very different from those of the ionized gas. From the inspection of Fig. \\ref{fig:greyvel}, the large scale velocity structure of the \\HMOL\\ shows two distinct velocity components.  One has a constant velocity gradient and corresponds to the extended rotating structure also seen in the ionized gas.  The second velocity component, only seen in \\HMOL\\ between 2\\arcsec\\ and 10\\arcsec, is reminiscent of a rotating ring with an inner radius of 6\\arcsec\\ and a keplerian fall-off. Fig. \\ref{fig:h2rot} shows that the \\HMOL\\ flux at NUC-1 within 2\\arcsec\\ has two secondary maxima at $\\simeq 1\\arcsec$ on both sides from the nucleus, as expected from a circumnuclear ring.  The same two maxima, however, are not observed at NUC-2. Comparing the projected radius at NUC-1 and the kinematical inner radius at NUC-2, the \"ring\" must be almost edge-on ($\\I>80\\arcdeg$), in agreement with the absence of the high velocity \\HMOL\\ components at NUC-1, OFF-1 and OFF-2 where the slit PA is perpendicular.  In summary, this putative molecular \"ring\" has an inner radius of $\\simeq 100\\PC$, is almost edge on ($\\I>80\\arcdeg$) and, from the velocity amplitude, includes a mass $M=0.6-2.5\\xten{9}\\Mo$. The uncertainty on the mass derives from the uncertainties on the inclination and the inner radius of ring.  This high velocity component is the \\HMOL\\ counterpart of the 100\\PC-scale molecular ring previously identified by several authors in CO and FIR \\citep{israel90,israel91,ryd93,hawa93}.  In CO, its major axis is oriented perpendicularly to the jet at PA$\\sim 145\\arcdeg$, its radius ranges between 80 and 140\\PC\\ and its rotational velocity ($\\simeq 220\\KM\\S\\1$ at $\\sim 120\\PC$, rescaling from  3\\MPC\\ to 3.5\\MPC\\ distance) results in an apparent dynamical enclosed mass of 1.5\\xten{9}\\Mo\\ \\citep{ryd93}.  A problem with the ring picture is that the \\HMOL\\ velocity decrease observed between 6\\arcsec\\ and 10\\arcsec\\ cannot be ascribed to circular keplerian motions.  The luminosity density plotted in Fig.  \\ref{fig:nuker}b and a Mass-to-Light ratio $\\Upsilon\\sim 1\\Mo/\\LKo$ naturally explain the mass enclosed by the \\HMOL\\ ring within $R=6\\arcsec$. However, the mass distribution plotted in Figure \\ref{fig:nuker}c is inconsistent with the observed velocity decrease between 6\\arcsec\\ and 10\\arcsec\\ from the nucleus, which would require negligible mass distributed in that range. Furthermore this model does not explain the presence of the other extended component.  The observed velocity field could be readily understood in terms of the gas dynamics of the gaseous bar observed by \\citet{mira99}.  We note that the two \\HMOL\\ velocity components form an approximate \"figure-of-8\" shape in the position-velocity diagram of Fig.  \\ref{fig:greyvel}. Similar velocity structures have been observed in edge-on spiral galaxies believed to contain barred potentials \\citep{merr99,bureau}. Hydrodynamical and N-body simulations \\citep{ab00,ba00} suggest that these characteristic velocity fields originate in weak barred potentials which support both the $x_1$ and $x_2$ families of orbits.  The slowly rotating component, seen in both ionized and molecular gas, is likely produced by the gas in the disk outside the bar, in combination with gas streaming along the $x_1$ family aligned with the axis of the bar. The \"anomalous\" \\HMOL\\ component is generated by gas within the Inner Lindblad Resonance (ILR) corresponding to the $x_2$ family of orbits. In most of the models, a gas pseudo-ring forms at or close to the location of the outer ILR (i.e. in our case the ring structure seen in \\HMOL\\ and CO). These models however do not include a BH. In the presence of a weakly barred potential and a supermassive BH, \\citet{wada} and \\citet{shlosh}  have shown that an additional nuclear Lindblad resonance (NLR) is formed close to the nucleus. Numerical results show that, on a dynamical time-scale, the gas is accumulated into a nuclear ring within the NLR \\citep{fu98}.  The gas ring becomes gravitationally unstable when the surface density of the gas reaches a critical value, fragments and forms a nuclear gas disk of several 10\\PC-radius \\citep{fu00}. This could provide a natural explanation for the rapidly rotating nuclear disk we see within the central 2\\arcsec\\ of Cen A, possibly accounting for the discrepancy between the apparent line of nodes observed in \\PA\\ and the kinematical one (end of Sec. \\ref{sec:fit}).  A problem for the bar picture could arise from the fact that the bar is not seen in the near infrared isophotes. As shown above, the \"ring\" system is close to edge-on and, depending on its actual shape, the barred potential might not be readily identified from the isophotes. In this case, it has been suggested that boxy or peanut-shaped bulges in edge-on systems can be associated to bars \\citep[e.g.][]{merr99,ba00}.  However, even in the near-IR, the analysis of Cen A isophotes is hindered by the high reddening and no definite conclusion can be drawn.  Another possibility is that the bar is a transient phenomenon, visible only in gas, as in the case of the Circinus galaxy \\citep{maio00}. This fact was suggested as a way to fuel an active galactic nucleus \\citep{shlosh}.  Though the explanation involving a gaseous bar is very appealing, nonetheless we cannot exclude that the \\HMOL\\ kinematics could be explained with a warped thin disk as in the models by NBT92 and Q92: in this case the velocity decrease of the ring-like \\HMOL\\ rotation curve at NUC-2 might be ascribed to the twisting of the disk line of nodes and/or changing inclination.  The \\HMOL\\ ring thus represents the inner edge of the dust lane and the inner radius of the ring might be identified with the tidal disruption radius of a molecular clouds in the nuclear gravitational potential.  The radius for tidal disruption of a spherical cloud in circular rotation around a spherical mass M is simply \\begin{equation} R_\\mathrm{tidal}=\\left(3\\frac{M}{m_\\mathrm{c}}\\right)^\\frac{1}{3}r_\\mathrm{c} = \\left(\\frac{9M}{4\\pi\\rho_\\mathrm{c}}\\right)^\\frac{1}{3} = 134\\PC \\left(\\frac{M_9}{n_4}\\right)^\\frac{1}{3} \\end{equation} where $m_\\mathrm{c}$, $r_\\mathrm{c}$ and $\\rho_\\mathrm{c}$ are the mass, radius and density of the cloud, $M_9$ is the mass in units of \\ten{9}\\Mo\\ and $n_4$ is the particle number density in units of \\ten{4}\\CM\\3. Since the mass enclosed within the ring is around $M_9\\simeq 1$, the tidal radius of a $\\simeq\\ten{4}\\CM\\3$ cloud is in excellent agreement with the observed inner radius of the ring. This molecular gas density, apart for being a reasonable value for a molecular cloud, is also similar to the value inferred by \\citet{israel90} from modeling the \\HMOL\\ emission.  In summary, even if the gaseous bar picture is very appealing in the framework of the fueling of the AGN, the warped disk explanation cannot be excluded and more data are required to prove the existence of a bar.  Summary}  New high spatial resolution, near-IR spectroscopy at several position angles and locations within the central 20\\arcsec\\ of the nearest radio galaxy Centaurus A have allowed us to derive velocities of ionized and molecular gas (\\FeII, \\PB, \\BG\\ and \\HMOL) that are not compromised by excitation effects or obscuration.  We have identified three distinct kinematical systems: (i) a \"nuclear\" rotating disk of ionized gas, confined within the inner 2\\arcsec, which is the counterpart of the \\PA\\ feature previously revealed by NICMOS imaging on the HST; (ii) a \"ring-like\" system with a $\\sim 6\\arcsec$ inner radius detected only in \\HMOL\\, the counterpart of the 100\\PC\\ scale \"torus\" detected in CO by other authors; (iii) a normal, \"extended\" component of gas rotating in the galactic potential.  The nuclear \"disk\", characterized by a low excitation spectrum, has an electron density in the range 3.5-4.5\\CM\\3\\ (in log) and its mass in ionized gas is $\\sim \\ten{3}\\Mo$. We have detected emission in the high ionization line \\SIX\\WL 12523.5\\AA; this is the first detection of a coronal line from an object characterized by a low excitation spectrum.  The nuclear rotation curves are well modeled by a thin disk in keplerian rotation around a central mass concentration. The line widths are consistent with unresolved rotation of the disk and there is no evidence for strong non-gravitational motions. The rotation curves suggest a central mass concentration that is dark ($M/L\\simeq 20\\,\\Mo/\\LKo$) and point-like at the spatial resolution of the data ($R<0\\farcs25\\sim 4\\PC$) which we interpret as a supermassive black hole. The minimum dynamical mass is $\\M>6\\xten{7}\\Mo$.  If we use the radio jet observations to exclude the case of a face-on nuclear disk ($\\I<15\\arcdeg$), we can constrain the mass to the range $\\M=2^{+3.0}_{-1.4}\\xten{8}\\Mo$.  The active nucleus in Centaurus A is accreting well below the Eddington limit, as found in other FR-I radio sources. A comparison between the X-ray luminosity and the FIR luminosity suggests that the AGN continuum lacks the big blue bump, a feature which, together with the low accretion efficiency, suggests an advection dominated accretion flow (ADAF).  The ratio of the BH mass to the galaxy mass is a factor of ten below the average of values directly measured in other galaxies suggesting either that it is overestimated or that it is characterized by a large spread. However, the BH mass over galaxy mass ratio of Centaurus A agrees with predictions derived from integrated quasar light and the same value measured for a sample of Seyfert 1 galaxies from reverberation mapping.  The ring-like system appears characterized by non-circular motions and a \"figure-of-8\" pattern observed in the \\HMOL\\ position-velocity diagram might provide kinematical evidence for the presence of a nuclear bar. However, with the present data, we cannot exclude that the same velocity systems could be explained with a warped-disk geometry.  Centaurus A is the nearest example of an active galaxy whose nucleus is heavily obscured by gas and dust.  Thus, despite its relative proximity, it has not been possible to use high resolution optical spectroscopy to study the kinematics of the gas near the putative black hole at the center of Cen A. The advent of high sensitivity, high spatial resolution near-IR spectroscopy on 8m-class telescopes has made it possible to resolve the multiple velocity components of the gas in the vicinity of the nucleus, and to obtain a dynamical BH mass estimate for this obscured nucleus.  Since dust obscuration is relatively common, and is likely to be even more so at higher redshifts, even higher spatial resolution near-IR spectroscopy, either from space or via adaptive optics on the ground, will be essential to extend this work to other obscured AGNs."
    },
    "0011/astro-ph0011435_arXiv.txt": {
        "abstract": "\\footnote{Contribution to the 6$^{th}$ Trieste Conference Proceedings, {\\it First steps in the origin of life in the Universe}, Trieste-Italy, 18-22 Sptember 2000} At early times the Universe was filled up with an extremely dense and hot gas. Due to the expansion it cooled below the binding energies of atoms which led to the formation of the first nuclei. In the physical environment of the post-recombination period of hydrogen, molecules such as $H_2$, $HD$ and $LiH$ can be formed. The proto-object formation, resulting from the growth of linear density fluctuations in the early Universe, can have an important impact on the chemical state of the Universe. Hence it can be enriched with metals, and thus lead to the formation of the first pre-biotic molecules. \\\\ In this contribution, I will present some scheme for the formation of primordial molecules and discuss the consequence of the formation of first stars on the existence of possible primordial pre-biotic. ",
        "introduction": "The conventional model of nucleosynthesis is the process by which chemical elements and their isotopes are formed, the heavy elements (carbon, oxygen, nitrogen and heavier ones) are thought to be the result of thermonuclear burning in stars, and especially the relatively rare stars which become supernovae \\cite{signore99}, \\cite{signore00}. \\\\ Standard big bang nucleosynthesis generated just a few elements: hydrogen, deuterium, helium and lithium, traces of beryllium and boron. The recombination of this primeval plasma led progressively to form the neutral species at $\\sim$ 10$^9$ yr after the big bang. Little is known about the history of the Universe corresponding to 10$^6$-10$^9$ years (or to redshifts of $z\\sim 1000-5$) after the big bang, and only recently astrophysicists begun to investigate seriously the post-recombination era and particularly the problem of the formation of the first molecules \\cite{lepp84}, \\cite{puy93}, \\cite{galli98}. \\\\ It is usually supposed that the first objects (first stars or galaxies) could be formed from a gas cloud by the gravitational attraction. The clouds contract adiabatically and some kind of {\\it cryogen} is necessary to form bound systems. Since the primeval gas does not contain heavy elements like carbon, primordial molecules such as $H_2$, $HD$ and $LiH$ are good candidates to lead to an efficient cooling \\cite{puy96}, \\cite{puy97}. In this context primordial molecules can play an important role in the pregalactic gas particularly on the thermochemical stability \\cite{puy98a}, \\cite{puy98b}. \\\\ In this contribution, I will present in Sect. 2 the formation of primordial molecules. Then in Sect. 3, I will describe how these molecules can influence on the formation of the first objects. In this context, I will discuss the important influence of these proto-objects on the second generation of molecular formation, which will be important for the formation of biomolecules. I will describe, in Sect. 4, some possible outlooks of this study. ",
        "conclusions": ""
    },
    "0011/astro-ph0011421_arXiv.txt": {
        "abstract": "We assess the feasibility of detecting star formation in damped Lyman-$\\alpha$ systems (DLAs) at $z>1$ through near-infrared spectroscopy using the forthcoming integral field units on 8\\,m-class telescopes. Although their relation to galaxies is not well established, high-$z$ DLAs contain most of the neutral gas in the Universe, and this reservoir is depleted with time -- presumably through star formation. Line emission should be an indicator of star formation activity, but searches based on Lyman-$\\alpha$ are unreliable because of the selective extinction of this resonant UV line. Using more robust lines such as H$\\alpha$ forces a move to the near-infrared at $z>1$. For line emission searches, spectroscopy is more sensitive than imaging, but previous long-slit spectroscopic searches have been hampered by the likelihood that any star forming region in the DLA galaxy disk would fall outside the narrow slit. The new integral field units such as CIRPASS on Gemini will cover sufficient solid angles to intercept these, even in the extreme case of large galactic disks at high redshift. On an 8\\,m-class telescope, star formation rates of $< 1\\,M_{\\odot}\\,{\\rm yr}^{-1}$ will be reached at $z\\approx 1.4$ with H$\\alpha$ in the $H$-band. Such star formation rates are well below $L^{*}$ for the high-$z$ Lyman-break population, and are comparable locally to the luminous giant H{\\scriptsize~II} complexes in M\\,101. It appears that instruments such as CIRPASS on Gemini will have both the sensitivity and the survey area to measure star formation rates in $z>1$ DLAs. These observations will probe the nature of damped Lyman-$\\alpha$ systems and address their relation to galaxies. \\index{abstract} ",
        "introduction": "Traditionally, flux-limited selection has been used to chart galaxy evolution. The measurement of quasar absorption lines allows an independent approach to studying the history of galaxies. The highest hydrogen column density absorbers seen in the spectra of background QSOs -- the damped Lyman-$\\alpha$ systems (DLAs) -- contain most of the neutral gas in the Universe at $z>1$ (ref.~1), although their exact relation to galaxies is not well established.  The global history of star formation can be inferred from the evolution in the co-moving density of neutral gas (predominantly in DLAs) as it is consumed in star formation (Fig.~1a \\& refs.~2,3).  The average star formation rate in each DLA depends then on their space density. One school of thought has $z>2$ DLAs being thick gaseous disks, the progenitors of massive spirals (e.g., ref.~4).  Alternatively, DLAs could be more numerous gas-rich dwarfs, potentially sub-galactic building blocks (e.g., Fig.~1b and refs.~5, 6\\,\\&\\,7).  Hence, to determine the nature of the DLAs requires a measurement of their star formation rate. This is most directly done through the hydrogen recombination lines (e.g., ref.~8), which should have luminosity proportional to the photo-ionizing flux produced by the most massive and shortest lived OB stars (and therefore trace the near-instantaneous star formation rate).  However, the largely unsuccessful searches for Ly$\\alpha$\\,121.6\\,nm emission do not provide useful limits as this line is resonantly scattered and hence is selectively suppressed by dust absorption, which is difficult or impossible to quantify.  Moving to the Balmer lines H$\\alpha$\\,656.3\\,nm and H$\\beta$\\,486.1\\,nm greatly reduces the problem of dust extinction, enabling the true star formation rates to be more accurately measured (e.g., ref.~9). For interesting redshifts where large samples of DLAs exist ($z>1$), these lines are redshifted into the near-infrared pass-bands.  \\begin{figure}[h] \\centering \\resizebox{0.49\\textwidth}{!}{\\includegraphics{bunkerf1.eps} } \\resizebox{0.49\\textwidth}{!}{\\includegraphics*[75,310][505,725]{bunkerf2.eps}\\hfil \\put(-140,100){\\scalebox{0.4}{\\includegraphics{bunkerf3.eps}}} } \\put(-300,143){\\large\\bf (a)} \\put(-130,142){\\large\\bf (b)} \\caption{{\\bf (a)} The observational limits on the star formation rates in high-reshifted damped Lyman-$\\alpha$ systems, based on H$\\alpha$ near-infrared long-slit spectroscopy with CGS\\,4 on UKIRT (ref.~10). The solid curve is the predicted average star formation rates for a closed-box model (ref.~3) for the hypothesis where spirals today evolve from high-redshift large disks.  The dashed curve is the prediction if present day spirals are assembled from hierarchical mergers of smaller, but more numerous, high-redshift DLAs (the hierarchical hypothesis). The curves plotted take no account of the possibility that the regions of star formation fall off the slit (which may be $>90$\\% of the cases using a spectroscopic slit).  Also shown are the predicted sensitivities of the CIRPASS instrument on Gemini, which will test {\\em all} these models, unlike the existing observations. The large area covered by the CIRPASS integral field unit means a much greater likelihood of intercepting any star forming regions. \\newline {\\bf (b)} A semi-analytic model of a halo with $V_{\\rm circ}=160$\\,km\\,s$^{-1}$ at $z=3$ (ref.~6).  The shaded region shows the area of the disk with column density in excess of $2\\times 10^{24}\\,{\\rm m}^{-2}$ (corresponding to DLAs). The overlaid CIRPASS fibres (0.25\\,arcsec across) shows that a single pointing of the integral field unit centred on a DLA will intercept most of the high column density gas -- the potential sites of star formation} \\label{fig:SFRs} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011084_arXiv.txt": {
        "abstract": " ",
        "introduction": "The ESA satellite Planck will survey the microwave sky with unprecedented sensitivity from 30 to 850 GHz. These observations will generate a large data set, with inhomogeneous data types, which will need to be accessible to users spread across many institutes in Europe and the US. How well it will be possible to mine this wealth of information depends critically on the proper design of the Data Management Component (DMC); that is on the efficiency and user-friendliness of the suite of tools that will be used to store, retrieve, process and query the data.  The IDIS \\cite{idis} project is a collaboration among the two Planck Data Processing Centres -- in particular OAT of Trieste, the Max-Planck Institut f{\\\"u}r Astrophysik and the Astrophysics Division of ESTEC. Within IDIS, we have started to develop the Planck DMC according to a 3-tier architecture. ",
        "conclusions": "This initial version of the Planck DMC allows maps from large sky surveys to be handled efficiently and it can be used to perform scientific data analysis. From the point of view of the client (scientist) who is interested in data processing: \\begin{itemize}  \\item the tool with its 3-tier design provides a convenient way to access the data. It allows the user to construct and handle ``scientific objects'' (such as maps or spectra) without having to understand the technicalities of data storage (formatting, optimization issues) \\item switching from a storage implementation which uses a file system on disk (e.g. a set of FITS files) or an object oriented database on a server is effortless \\item the use of a data structure such as HEALPix has proved to be a crucial factor for the speed of operations involving Spherical Harmonic decomposition \\end{itemize} From the point of view of the software engineer who develops the data storage system: \\begin{itemize}  \\item a database offers a way of handling and managing the data which is more powerful than a simple file system \\item from benchmarking tests, the use of an object oriented database in our case proved to be preferable to a relational database: for speed of data access and the option of storing objects of any given complexity (that is objects which refer to other objects, which in turn refer to other objects and so forth) \\end{itemize}  Using this version of the Planck DMC we have derived the spectral indices of radio diffuse emission from the available large-sky surveys. At high Galactic latitudes these are consistent with the spectral indices derived previously from the Haslam map and the Reich \\& Reich map (\\cite{Tegmark96}; \\cite{bouchet96}; \\cite{Bouchet99}), but are more precise (Giardino et al., in preparation)."
    },
    "0011/astro-ph0011567_arXiv.txt": {
        "abstract": "The methodology involved in deriving the Hubble Constant via the calibration of the corrected peak luminosities of Type~Ia supernovae (SNe) is reviewed.  We first present a re-analysis of the Cal\\'an-Tololo (C-T) and Center for Astrophysics (CfA) Type~Ia SN surveys. Bivariate linear least squares and quadratic boot-strapped fits in peak apparent magnitude and light curve shape are employed to correct this heterogeneous sample of peak apparent magnitudes, resulting in an homogeneous (and excellent) secondary distance indicator: the so-called corrected peak luminosity. We next provide an empirical calibration for this corrected luminosity, using Cepheid-based distances for seven nearby spiral galaxies host to Type~Ia SNe.  Included in this sample is the spectroscopically peculiar SN~1991T (in NGC~4527), whose corrected peak luminosity is shown to be indistinguishable from that of so-called ``normal'' SNe.  A robust value of the Hubble Constant is derived and shown to be H$_0$=73$\\pm$2(r)$\\pm$7(s)\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "introduction": "The maximum brightness of a Type~Ia supernova (SN) is \\it not \\rm a standard candle.  As illustrated graphically by Hamuy et~al. (1996) and Riess et~al. (1996), there is an order of magnitude span seen in the luminosities at the peak of the SNe light curves in their respective samples (depending upon the bandpass in question).  On the other hand, an outstanding standard candle \\it can \\rm be constructed by applying a light curve shape-dependent correction factor to the observed peak magnitude (e.g., Phillips et~al. 1999; Jha et~al. 1999; Suntzeff et~al. 1999; Gibson et~al. 2000; Freedman et~al. 2001; Gibson \\& Stetson 2001). This corrected peak luminosity succeeds as a standard candle where others fail because (i) the extreme luminosity of Type~Ia SNe allows one to probe substantially further into the unperturbed Hubble flow (in comparison with competing secondary indicators such as the Tully-Fisher relation or surface brightness fluctuations), and (ii) the intrinsic scatter is smaller than that of other indicators.  Calibrating this secondary distance indicator can be approached from a purely theoretical tack (e.g., H\\\"oflich \\& Khokhlov 1996), but the numerous model uncertainties still lead the conservative to favour the empirical calibration. The \\it Sandage/Tammann/Saha Type~Ia SN HST Calibration Project \\rm (Saha et~al. 1999) was designed to provide just such a calibration, by determining the Cepheid distances to eight nearby galaxies (IC~4182, NGC~5253, 4536, 4496A, 4639, 3627, 3982, 4527) host to Type~Ia SNe.  This original program has been supplemented with two additional calibrators situated within NGC~3368 (Tanvir et~al. 1999) and NGC~4414 (Turner et~al. 1998).  Table~1 of Gibson et~al. (2000) provides a complete list of the galaxies and SNe covered by all these calibration programs, including a subjective quality ranking.  In what follows, we review the steps involved in calibrating the Type~Ia SNe extragalactic distance scale.  We first provide a consistent set of linear and quadratic fits to the peak magnitude-light curve shape data from the Cal\\'an-Tololo (Hamuy et~al. 1996) and Centre for Astrophysics (Riess et~al. 1998,1999) SNe samples.  These fits are then shown to provide a robust measure of the Hubble Constant when the relations are calibrated using the seven best (currently available) nearby Type~Ia SNe with Cepheid-based host galaxy distance determinations. ",
        "conclusions": "A re-analysis of the combined Cal\\'an-Tololo and Center for Astrophysics Type~Ia SNe datasets is presented, with new linear and quadratic fits to the empirical peak magnitude-decline rate relationship determined. The fit coefficients are extremely robust to uncertainties in photometry, reddening, and peculiar velocity.  With a calibration to the corrected Hubble Diagrams now provided by seven high-quality nearby calibrators with accurate host galaxy Cepheid distance determinations, our favoured value for the Hubble Constant becomes: \\begin{center} \\fbox{${\\rm H}_0 = 73\\pm 2({\\rm r})\\pm 7({\\rm s})\\;{\\rm km\\,s^{-1}\\,Mpc^{-1}}$} \\end{center} \\noindent The systematic uncertainties clearly limit the precision to which we can determine H$_0$; until significant progress is made in improving our understanding of, for example, the distance to the LMC and the sensitivity of the Cepheid PL relation to metallicity, the uncertainty associated with Cepheid-based Hubble Constant \\it cannot \\rm be reduced below 10\\%."
    },
    "0011/astro-ph0011201_arXiv.txt": {
        "abstract": "Many recent estimates of the star formation rate density at high redshift rely on rest-frame ultraviolet (UV) data.  These are highly sensitive to dust absorption.  Applying a correlation between the far-infrared (FIR) to UV flux ratio and UV color found in a local starbursts to galaxy samples out to $z \\sim 3$, one can account for most of the FIR background.  However, the correlation is based on a sample that does not include the most extreme starbursts, Ultra Luminous Infrared Galaxies (ULIGs).  Our new UV images of ULIGs show that their FIR fluxes are underpredicted by this correlation by factors ranging from 7 to 70.  We discuss how ULIGs compare to the various types of high-$z$ galaxies: sub-mm sources, Lyman Break Galaxies, and Extremely Red Objects. ",
        "introduction": "About 60\\%\\ of the {\\em intrinsic} bolometric luminosity of star forming populations is emitted between 912\\AA\\ and 3000\\AA\\ with little variation in this fraction with star formation duration (from models of Leitherer \\etal\\ 1999).  Hence UV light is potentially very useful for measuring star formation.  It represents direct emission from hot main-sequence stars, the same stars that will provide the majority of the mechanical energy feedback into the ISM.  The UV spectrum is very rich in features that can be used to diagnose the stellar populations and intervening ISM.  The overall intrinsic spectral slope in the UV is fairly constant.  For young ionizing populations it is set by the Rayleigh-Jeans tail of the Planck function.  When looking at high redshifts the importance of the rest-frame UV increases; at a fixed bandpass it is the last stop before the Lyman-edge.  The big problem with using UV measurements is dust, which most efficiently absorbs and scatters UV radiation.  The absorbed UV light is reradiated at far infrared (FIR) wavelengths.  Not only is the amount of dust important, so to is its distribution.  If the geometry is unfavorable, as in the mixed stars and dust model, then the UV emission will be dominated by the stars closest to the observer and the center may be practically invisible (e.g.\\ Witt, Thronson, \\&\\ Capuano 1992).  Fortunately, this is not the case for a wide range of starburst galaxies as shown in Fig.~1.  For them the FIR to UV flux ratio, or infrared excess (IRX), which gives the effective dust absorption, correlates with the UV spectral slope $\\beta$ (Meurer \\etal\\ 1995, 1999).  This is a prediction of the dust screen model (Witt \\etal).  While variations on this geometry can also explain this correlation (Calzetti 1997; Charlot \\&\\ Fall 2000), the mere fact of this correlation allows one to recover the intrinsic luminosity of starbursts using UV quantities alone. ",
        "conclusions": ""
    },
    "0011/astro-ph0011037_arXiv.txt": {
        "abstract": "The population of the Kuiper Belt within 50~AU of the Sun has likely been severely depleted by gravitational perturbations from the giant planets, particularly Neptune.  The density of Kuiper Belt objects is expected to be two orders of magnitude higher just beyond 50 AU, where planetary perturbations are insignificant.  In 1998 and 1999, we surveyed for Kuiper Belt Objects (KBOs) in 6 fields of the ecliptic (total sky area 1.5 deg$^2$) to limiting magnitudes between $R=24.9$ and $R=25.9$.  This is deep enough to detect KBOs of diameter $\\gtrsim 160$~km at a distance of 65~AU.  We detected 24 objects.  None of these objects, however, is beyond 53~AU.  Our survey places a 95\\% CL upper limit of $\\Sigma < 5\\,{\\rm deg}^{-2}$ on the surface density of KBOs larger than $\\sim160$~km beyond 55~AU.  This can be compared to the surface density of $\\sim6\\,{\\rm deg}^{-2}$ of $\\ge160$~km KBOs at distances 30--50~AU determined from this survey and previous shallower surveys.  The mean volume density of $D>160$~km KBOs in the 55--65~AU region is, at $>95\\%$ confidence, less than the mean density in the 30--50~AU region, and at most \\twothirds\\ of the mean density from 40--50~AU.  Thus, a substantial density increase beyond 50~AU is excluded in this model-independent estimate, implying that some process or event in the history of the Solar System has truncated the distribution of 160-km planetesimals at $\\sim50$~AU.  A dense primordial disk could be present beyond 50~AU if it contains only smaller objects, or is sufficiently thin and inclined to have escaped detection in our 6 survey fields. ",
        "introduction": "Our planetary system should not be expected to be bounded by the orbit of the ninth planet Pluto, as pointed out by \\citet{Ed49} and \\citet{Ku51}.  Our current general understanding of the formation of our solar system is consistent with the expectation that a large population of small bodies -- leftovers from the primordial planetesimal disk in the Solar Nebula -- exists at the present time beyond the orbits of Neptune and Pluto where planetary accretion timescales exceed the current age of the solar system.  Indirect evidence for such a population (now known as the {\\it Kuiper Belt}) exists in the orbital properties of the short period comets, as has been suggested by several authors \\citep{Fe80,Du88}. Direct observational evidence was first obtained with the discovery of 1992 QB1 \\citep{Je93}, and has subsequently grown with the detection of more than 300 objects.  It is notable that all but one of the currently known Kuiper Belt objects (KBOs) lie within 55 AU of the Sun.  The semimajor axis zone $30<a<50$ AU is dynamically a very complex region. The primordial Kuiper Belt population here is expected to have been extensively sculpted and depleted throughout solar system history [see review by \\citet{Ma00}].  Numerical simulations show that most primordial KBOs on orbits with $a<36$~AU would have been ejected from the Kuiper Belt in much less than the age of the solar system \\citep{Ho93, Le93, Du95}.  In the $36<a<42$ AU zone, the dynamical lifetime of low eccentricity, low inclination orbits is comparable to the age of the solar system, whereas beyond $a\\approx42$ AU the dynamical lifetime is well in excess of the age of the Solar system. The orbital migration of the giant planets in early solar system history is expected to have rearranged the primordial orbital distribution in the $a<50$ AU zone, by sweeping a significant fraction of KBOs into eccentric orbits at narrow semimajor axis zones at the locations of mean motion resonances with Neptune \\citep{Ma93, Ma95}. Furthermore, there exists a significant ``scattered'' KBO population, in very eccentric orbits with perihelia near $\\sim 36$ AU; these KBOs are thought to have been formed in the $36< a< 42$ AU zone and been gravitationally scattered during a close encounter with Neptune \\citep{Le97}.  The current estimate of the surface density in the inner Kuiper Belt is approximately two orders of magnitude less than that expected from a smooth extrapolation of the surface density from the planetary region [{\\it cf.} \\citet{We77}].  It is also approximately two orders of magnitude less than that required for the formation of the largest observed KBOs within $10^8$~yrs, before perturbations from the giant planets curtailed their growth \\citep{S96, SCb, KL99}.  These arguments suggest severe depletion of the Kuiper Belt within 50 AU.  Beyond 50 AU, the gravitational influence of Neptune and the giant planets is insignificant.  A higher present-day density of objects, reflective of the primordial surface density of solids in the Solar Nebula, would be expected in this unperturbed distant Kuiper Belt. Previous surveys have detected no objects beyond 50~AU, leading several authors \\citep{Do97, Ch99, JLT} to suggest that this higher-density outer region is in fact absent, and that the primordial planetesimal disk must have had a cutoff near 50~AU. \\citet{Gl98} calculate the fraction of detected KBOs expected to be beyond 50~AU; applying their formula to the number of KBOs known today also indicates that there is a deficiency of distant objects.  These conclusions, however, depend quite sensitively upon the exponents of assumed power laws for size and radial distributions of the population.  In this Letter we report the results of a survey for faint KBOs covering more than one square degree of the ecliptic.  Most of the known KBO population lie at $\\sim40$~AU and have magnitudes $R\\le24$. At our limiting magnitudes of $R\\ge25$, we can detect physically similar objects to a distance of 65~AU or greater.  We compare the results of our faint survey with the detected population of similarly sized KBOs at smaller distances, and test whether the population density increases beyond 50 AU.  This comparison of the population of $\\gtrsim 160$~km objects is independent of any models or assumptions for the size distribution of the KBOs in the two regions. ",
        "conclusions": "If planetary perturbations are solely responsible for the structure of the Kuiper Belt, a dense primordial disk would be expected beyond $\\sim50$~AU where these perturbations are insignificant.  Our survey could have detected such a disk but did not.  There are several possible explanations for this non-detection.  \\begin{enumerate}  \\item KBOs in the outer Kuiper Belt could be fainter than expected, owing to a lower albedo or much redder color or much smaller sizes. To explain our survey results, this would require a change in the physical properties of KBOs beyond 50~AU.  \\item The outer Kuiper Belt could have been dynamically excited early in the history of the solar system by a stellar encounter \\citep{Id00} most likely during the Sun's residence in its birth cluster \\citep{AL00}; by perturbations from large Neptune-scattered planetesimals \\citep{P99} or by proto-planetary cores \\citep{Th99}.  Such excitation could have increased the orbital inclinations and eccentricities of the objects, lowering the apparent surface density on the sky and possibly decreasing the maximum size of objects through the cessation or reversal of the accretion process.  More extreme excitation could have stripped away the outer Kuiper Belt.  \\item The disk is actually present, but dynamically cold \\citep{Ha00}. It could have escaped our survey by simply not intersecting our fields. This is entirely possible, if this disk is inclined to the ecliptic plane by as little as 1\\arcdeg. The invariable plane (the angular momentum plane of the solar system), inclined to the ecliptic by 1.5\\arcdeg, would be a likely candidate for the location of such a cold, dense disk. Our Field~G is only 0\\fdg6 from the invariable plane, so the scale height of the cold disk would have to be $\\ll1\\arcdeg$ to have escaped detection.  This implies a very thin, extremely dense disk. Absence of $>50~AU$ detections in future deep survey fields would eliminate the possibility of a cold disk.  \\end{enumerate}"
    },
    "0011/astro-ph0011347_arXiv.txt": {
        "abstract": "We combine AAVSO and VSS/RASNZ optical and {\\it Extreme Ultraviolet Explorer\\/} EUV light curves of dwarf novae in outburst to place constraints on the nature of dwarf nova outbursts. From the observed optical--EUV time delays of $\\approx 0.75$--1.5 days, we show that the propagation velocity of the dwarf nova instability heating wave is $\\approx 3~\\rm km~s^{-1}$. ",
        "introduction": "Dwarf nova outbursts are understood to be the result of an instability in the rate of mass transfer through the accretion disk surrounding the white dwarf in these semidetached binaries. The instability can be triggered at large or small disk radii, resulting in normal, fast-rise outbursts or anomalous, slow-rise outbursts, respectively. In either case, the beginning of the outburst is signaled by a rise of the optical flux, followed by a rise of the UV flux as material sinks through the disk, converting gravitational potential energy into rotational kinetic energy and radiation. This is followed by a rise in the EUV flux as material passes through the boundary layer between the disk and the surface of the white dwarf, where it converts its prodigious rotational kinetic energy into radiation.  Because the optical and EUV flux of dwarf novae is produced in physically distinct regions of the accretion disk, light curves in these wavebands provide important diagnostics of the nature of dwarf nova outbursts. Toward this end, we present optical and EUV light curves of dwarf novae observed by the {\\it Extreme Ultraviolet Explorer\\/} ({\\it EUVE\\/}) satellite. ",
        "conclusions": "The anomalous outbursts of SS~Cyg (Fig.~1) manifest the gradual increase of the optical and EUV light curves expected for inside-out outbursts. The optical and EUV flux rises during the beginning of these outbursts as the heating wave sweeps outward through the disk, causing more and more material to flow through the disk and boundary layer onto the white dwarf. In contrast, the normal outbursts of SS~Cyg (Fig.~2), U~Gem (Fig.~3), and VW~Hyi (Fig.~4) manifest the fast increase of, and the delay between, the optical and EUV light curves expected for outside-in outbursts. As measured from the initial rise of the optical light curves, the delay of the rise of the EUV light curves is $\\approx 1.5$, 1.25, and 0.75 days for SS~Cyg, U~Gem, and VW~Hyi, respectively. In VW~Hyi, the FUV light curve rises $\\approx 0.5$ days after the optical light curve and $\\approx 0.25$ days before the EUV light curve, but falls as slowly as the optical light curve, consistent with the expectation that the accretion disk and not the boundary layer is the source of the FUV flux. We were not able to observe the rise of the EUV light curves of the superoutbursts of VW~Hyi (Fig.~5) or OY~Car (Fig.~6), but during both observations of OY~Car the EUV light curve was observed to fall and then rise while the optical light curve was declining only slowly. This behavior is evidence for a decrease and subsequent increase of the mass-accretion rate onto the white dwarf, as might be expected if a normal outburst is rejuvenated by an increase in the mass-accretion rate driven by the tidal instability expected for such high mass-ratio binaries.  The observed optical-EUV delays of the normal outbursts of SS~Cyg, U~Gem, and VW~Hyi provide the most direct measurement of the velocity of the heating wave which transforms the disk from quiescence to outburst. Assuming the system parameters shown in Table~1, that the radius of the disk is $R_{\\rm disk}\\approx 0.7\\times R_{\\rm L1}$, and that the disk instability starts at the outer edge of the disk, the velocity of the heating wave $v\\approx R_{\\rm disk}/{\\rm delay}\\approx 3~\\rm km~s^{-1}$. This result is consistent with $v=\\alpha c_{\\rm s}$ if the viscosity parameter $\\alpha\\approx 0.2$ and the sound speed $c_{\\rm s} =10\\, (T/10^4~{\\rm K})^{1/2}\\approx 15~\\rm km~s^{-1}$.  \\begin{table} \\vskip -0.2cm \\caption{System Parameters$^a$ and Velocity of the Heating Wave} \\begin{tabular}{llll} \\tableline Parameter&\tSS Cyg&\tU Gem&\tVW Hyi\\\\ \\tableline \\hbox to 2.5in{$P_{\\rm orb}$ (days)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 0.2751&\t0.1769&\t0.0743\\\\ \\hbox to 2.5in{$M_1$ ($M_{\\sun}$)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 1.19&\t  1.26&\t  0.63\\\\ \\hbox to 2.5in{$M_2$ ($M_{\\sun}$)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 0.70&\t  0.57&\t  0.11\\\\ \\hbox to 2.5in{$q = M_1/M_2$\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 1.69&\t  2.17&\t     6\\\\ \\hbox to 2.5in{$R_{\\rm L1}/a$\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 0.43&\t  0.45&\t  0.54\\\\ \\hbox to 2.5in{$a = [(P_{\\rm orb}/2\\pi )2G(M_1+M_2)]^{1/3}$ (cm)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& $1.5\\times 10^{11}$&\t$1.1\\times 10^{11}$&\t$4.7\\times 10^{10}$\\\\ \\hbox to 2.5in{$R_{\\rm disk}\\approx 0.7\\times R_{\\rm L1}$ (cm)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& $4.6\\times 10^{10}$&\t$3.5\\times 10^{10}$&\t$1.8\\times 10^{10}$\\\\ \\hbox to 2.5in{delay (days)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}& 1.5&\t1.25&\t0.75\\\\ \\hbox to 2.5in{$v\\approx R_{\\rm disk}$/delay ($\\rm km~s^{-1}$)\\leaders\\hbox to 5pt{\\hss.\\hss}\\hfil}&\t3.5&\t3.3&\t2.7\\\\ \\tableline \\end{tabular} \\phantom{a}$^a$Ritter, H., \\& Kolb, U. 1998, \\aap , 129, 83 \\end{table}"
    },
    "0011/astro-ph0011492_arXiv.txt": {
        "abstract": "We have identified a large sample  of probable low mass members of the young open cluster IC2391 based on optical (VRIZ) and Infrared (JHK$_s$) photometry. Our sample includes 50 probable members and 82 possible members, both very low mass stars and brown dwarfs. We also provide accurate positions for these stars and brown dwarf candidates  derived from red UK Schmidt plates measured using the microdensitometer SuperCOSMOS. Assuming an age of 53 Myr, we estimate that we have reached a mass of 0.025 M$_\\odot$, if the identified objects are indeed members of IC2391. ",
        "introduction": "IC2391 is one of  the youngest and nearest open clusters. The  {\\it Hipparcos} distance modulus for the cluster  is (m-M)$_o$$\\sim$5.82. As a compromise between various determinations of the distance modulus and {\\it Hipparcos} parallaxes (Becker \\& Fenkart 1971; Lyng\\aa{ } 1987; van Leeuwen 1999; Robichon et al.  1999),  we adopt (m-M)$_0$=5.95$\\pm$0.1, corresponding to a distance of 155 pc. The mean reddening towards the cluster is small, E(B-V)=0.06 (Patten \\& Simon 1996). Prior to our lithium depletion age determination, the most commonly quoted age estimate for IC2391 was $\\sim$35 Myr (Mermilliod 1981). In Barrado y Navascu\\'es, Stauffer \\& Patten (1999) we derived a new age estimate of $53 \\pm 5$ Myr.  Because of its proximity and youth, IC2391  has been the target of a number of recent studies. Rotational velocities, lithium abundances and H$\\alpha$ data for G, K and early M  dwarf members can be found in Stauffer et al. (1989, 1997). X-ray fluxes, rotational periods and low resolution spectra for another set of probable cluster members are provided by Patten \\& Simon (1986) and Simon \\& Patten (1998). All these data support the conclusion that IC2391 is  younger than the Alpha Per cluster ($\\sim$85 Myr in the new lithium age scale, Stauffer et al. 1999; Barrado y Navascu\\'es, Stauffer, \\& Bouvier 1999; Stauffer \\& Barrado y Navascu\\'es 2000) but significantly older than 10 Myr (e.g., because the cluster contains no classical TTauri stars).  For an age of 35 Myr, we expected the lithium depletion boundary for IC2391 to be at I$_C$$\\le$15.2 (Baraffe 1998, private communication; Stauffer, Schultz, \\& Kirkpatrick 1998). The ROSAT images of the cluster only supplied candidate cluster members down to I$_C$$\\sim$15. No appropriate deep photographic plate sequences are available for a proper motion survey to the desired depth. We therefore obtained deep, multicolor photometric imaging of the cluster in order to produce a list of candidate cluster members. Because of the low galactic latitude of the cluster ({\\it b}$^{\\sc II}$=-6.90), our derived candidate list may be significantly contaminated by field stars. We therefore defer extensive discussion of these candidates until detailed spectroscopic follow-up is available.  In Section 2,  we  provide the details of our imaging program. The method we used to identify candidate cluster members is outlined in Section 3. ",
        "conclusions": "\\subsection{Initial selection of candidate members of IC2391}  The initial selection of candidate members of IC2391 was carried out using the location of the detected stars on the I$_C$ versus (R-I)$_C$ color-magnitude diagram. Figure 2 displays  the data obtained at the 4m CTIO/BTC telescope, where detections are shown as dots. An empirical Zero Age Main Sequence (ZAMS), based on data from Leggett (1992),  is included as a solid line.   The ZAMS is plotted for our assumed IC2391 distance modulus and a reddening of E(R-I)=0.007   [David:  is this E(R-I) compatible with the E(B-V) mentioned earlier in the text?] Completeness and limiting magnitudes are depicted as dashed lines. We selected all the detected objects in a wide band well  above the ZAMS. The lower envelope for the strip follows a 50 Myr isochrone (D'Antona \\& Mazzitelli 1997) and the upper envelope is displaced 0.75 mag brighter to include equal mass binaries. The boundaries of the strip were also adjusted to take into account the   photometric errors and errors in the distance, age  and the reddening. Hence, this band is  wider than the 0.75 mag. Then, all candidates  were visually inspected on the original images, in order to avoid the presence of false detections, and  to verify the stellar-like shape of the detections and the lack of nearby bright stars or cosmic rays which could modify the photometry of the candidate. Since there is no clear separation between the field stars and the location of the cluster isochrone, we expect a strong contamination by spurious members. Specifically, those candidates located only slightly above the ZAMS should be considered with some caution, since they are well below the 50 Myr isochrone. In total, we have selected 206 candidate members, 94 from the BTC survey and another 112 from the 0.9m and 1.5m data. Figure 3 shows all our candidate members as circles. Solid circles indicate the position of the candidates whose membership has been confirmed  spectroscopically, via radial velocity, spectral type and several spectroscopic features (H$\\alpha$, NaI8200\\AA). These results are discussed in Barrado y Navascu\\'es, Stauffer,  \\& Patten  (1999), where we provide a lithium depletion boundary (see Barrado y Navascu\\'es, Stauffer \\& Bouvier 1999) age for the cluster (53$\\pm$5 Myr). Figure 3 also displays the position of a 50 Myr isochrone (short-dashed line), adapted from D'Antona \\& Mazzitelli (1997). Based on our spectroscopically confirmed members and on Simon \\& Patten (1996) data, we have created an empirical IC2391 isochrone, shown as a long-dashed line.  Table 2 lists the positions, optical and infrared photometry, the separation  between the optical detection and the IR source and the identification with stars from the literature, for  each candidate.   \\subsection{Color-Magnitude and Color-Color Diagrams of IC2391}  The merger of our RI$_C$ data with the 2MASS  photometry, and the additional VI$_C$Z photometry collected in January 1999, allows us to create a large database of optical-infrared broadband photometry in 7 different filters. Therefore, we have been able to construct several color-magnitude and color-color diagrams. We have used these diagrams to estimate the membership status of each candidate.  Figure 4a depicts the [V,(V-I$_C$)] diagram. The empirical main sequence (MS) for young disk stars by Leggett (1992) is included as a solid line. This MS was built for M0-M9 dwarfs, in the ranges 6.65 $\\le$ M(I$_C$) $\\le$ 14.67, 4.77 $\\le$ M(K$_{cit}$) $\\le$ 10.17, 0.75 $\\le$ (R-I)$_C$ $\\le$ 2.30, and 0.17 $\\le$ (H-K) $\\le$ 0.48. Note that our 2MASS data were taken in K$_s$. However, there are  no important differences between  the K$_s$ and K$_{cit}$\\ systems (see Persson et al 1998). For comparison purposes, the shift of the photometry corresponding to an interstellar absorption of A$_v$=2 is plotted as an arrow. Proposed members of IC2391 from the literature (Patten and Simon 1996; Patten and Pavlovski 1999) appear as crosses. Figure 4b displays the I$_C$ magnitudes against the (I$_C$-K$_s$) color, whereas  Figure 4c shows the K$_s$ magnitudes against the (J-K$_s$). In all these figures,  solid triangles represent initial candidate members whose membership has been rejected  based on  Figures 4 and 5 (probable non-members), whereas open triangles  show the possible non--members. Possible and probable members are represented by open and solid circles, respectively (see next subsection).    Several color-color diagrams can be found in Figure 5. Symbols are as in Figure 4. Panel a and b depict [(V-R$_C$),(R-I)$_C$], and  [(I$_C$-K$_s$),(R-I)$_C$], respectively. It is clear from examination of these diagrams that our original sample included both stars that are plausible cluster members and objects that are instead very likely to be heavily reddened, background stars.   \\subsection{A final list of photometric candidate members}  In order to  remove the spurious members present in our candidate member list, we have used several color-color diagrams  and color-magnitude diagrams. This selection was carried out in a hierarchical way, stressing the spectroscopic and IR data. The scheme we followed is described below:  \\begin{enumerate}  \\item A fraction of our candidates has intermediate resolution spectroscopy (Barrado y Navascu\\'es et al. 1999), which indicates whether they are probable members or non-members.  \\item We have removed the objects having a large interstellar reddening (see Figure 4b,c and Figure 5b). They were classified as probable non-members.  \\item Color-color diagrams  were used to remove additional  probable non-members (Figures 5a,b).  \\item  Objects fainter than I$_C$=17 mag, and without IR data (they are too faint to be detected by the 2MASS survey) were classified as possible members, whereas stars brighter than this value, with no IR data (they should have been detected by 2MASS) are listed as possible non-members.  \\item Objects bluer than the Leggett's (1992) IR  main-sequence appear in Table 2 as possible members.  \\item Objects bluer than our empirical optical  isochrone or  at the upper (bright) edge of our CM diagram selection strip were also classified as possible members.  \\item Finally, the  objects that remain are considered to be bona fide members of the cluster (``probable members\").  \\end{enumerate}   As a summary, we have classified our initial IC2391 candidate members in four different categories:  \\begin{itemize}  \\item  Probable members. Objects located in all  CM and CC diagrams with positions which indicate membership. They are identified with the flag ``{\\it MEM}'' in the last column of Table 2 (50 objects, including 16 having spectroscopy).  \\item Possible members (82 objects). They appear identified with the flag ``{\\it MEM}?'' in Table 2.  \\item Possible non-members (10). (``{\\it NM?}'' flag).  \\item Probable non-members (64). All of them  are flagged with ``{\\it NM}'' in  Table 2.  \\end{itemize}  If we only consider our probable  members, we have detected objects $\\sim$2 magnitudes fainter than the previous surveys (Rolleston \\& Byrne 1997; Simon \\& Patten 1998; Patten \\& Pavlovsky 1999). Our faintest possible IC2391 member has I$_C$=20.9, which is 5 magnitudes fainter than the least massive candidate discovered to date. In total, we list 132 objects as probable or possible members of IC2391.   We have  not  computed the luminosity function (LF) and mass function (MF). Due to its low galactic latitude  of the cluster ({\\it b}$^{\\sc II}$=-6.90), even when using the probable/possible member list, pollution by spurious members is likely to be significant. Visual inspection of the I,(R-I)$_C$ color-magnitude diagram (Figure 3) reveals two important characteristics:  the relative high number of candidate members with I$_C$ in the range 19.5--20.9, and the apparent gap just before this clustering occurs. only additional photometry or spectroscopy would allow us to establish if these faint candidates are real. However, we expect a strong contamination by field stars in  this range. If most of these objects are, indeed, not members, the gap would be illusory and the low number of candidate members below I$_C$=18  could be a consequence of our completeness limit (dotted line in Figure 3). The question is still open, until a follow-up spectroscopic study is carried out.  \\subsection{The contamination by field stars}  In a deep optical survey of the Pleiades, Bouvier et al. (1998) estimated that the contamination due to field stars was 25\\%. This value has been confirmed by subsequent spectroscopic follow-up of the Pleiades candidates. Since  IC2391 is closer to the Galactic Plane than the Pleiades ({\\it b}$^{\\sc II}$=-6.90 and  {\\it b}$^{\\sc II}$=-23.52, respectively), the contamination should be much stronger. In fact, Figure 2 of Bouvier et al. (1998) shows a clear   discontinuity between field stars and the Pleiades population.  This is not the case for IC2391, whose color-magnitude digram (Figure 3) depicts  a smooth transition between the field and  the locus of the main sequence of the cluster, indicating that contamination by spurious members should be stronger than in the case of the Pleiades. To estimate the degree of contamination by field stars, we have constructed histograms of the number of stars per magnitude bin at various (R-I)$_C$ color intervals. In Figure 6 we show the histogram for the (R-I)$_C$=1.9-2.0 range, representative of the distribution of stars down to approximately the completeness limit of the BTC data. The locus of the cluster is indicated. Similar diagrams for adjoining color ranges provide analogous results. The last bin before the cluster contains 3 stars, whereas the brightest bin, well above the main sequence for equal mass binary  cluster stars (0.75 brighter than the single star MS), has 5 stars. The bins in between have 17 stars. Assuming  an average contamination of 4 stars per bin, the pollution rate would be $\\sim$50\\% for the initial list of IC2391 candidates (the RI$_C$ survey). Our optical and infrared color-color and color-magnitude diagrams have allowed us to remove about 33\\% of the candidates from  the initial 206 objects. Therefore,  we estimate that objects cataloged as  members in Table 2 (probable and possible members) still have a $\\sim$25\\% or greater probability of being spurious.  The faintest objects, lacking VZJHK data, may be more strongly polluted by field stars.       \\subsection{Brown Dwarfs in the Cluster}  Based on the position of the lithium depletion boundary in members of IC2391, Barrado y Navascu\\'es, Stauffer,  \\& Patten (1999) have recently determined an age for the cluster of 53$\\pm$5 Myr. For this age, the border between the stellar and substellar regimes would be at M(I$_C$)=11.06 (Baraffe  2000, private communication). Taking into account the  IC2391 distance and reddening, this leads to I$_C$=17.03. Therefore, all our IC2391 candidates fainter than that value should be brown dwarfs if they indeed belong to the cluster. Our list of final candidates contains ten objects which have been cataloged as probable and possible  members based on both optical and infrared data and  which have I$_C$ in the range 17.06--17.62. Using the Baraffe (2000) models, their mass range is 0.070--0.055 M$_\\odot$. Actually, two of these objects (CTIO-061 and CTIO-113) have been observed using   low S/N, intermediate resolution spectroscopy (Barrado y Navascu\\'es, Stauffer,  \\& Patten 1999). Several spectral characteristics, such as H$\\alpha$ and NaI8200\\AA{ } equivalent widths, spectral type, rough radial velocities, indicate that they are real members of the cluster and, therefore, brown dwarfs.  Another 50 objects  in Table 2 are fainter than than the I$_C$ magnitude of the stellar/substellar limit and lack infrared data. All of them  are listed as possible members in Table 2, with magnitudes down to I$_C$=20.9, M(I$_C$)=14.93, and M $\\ge$ 0.025 M$_{\\odot}$.    \\subsection{Comparison with previous surveys and  other clusters}  A comparison between our final candidate members  (circles) and previous surveys (triangles) of IC2391 is shown in Figure 7a. Probable candidate members are displayed with solid circles, whereas open circles represent objects which  do not have infrared counterparts (normally, because they are too faint to be detected by 2MASS). The stars from Simon \\& Patten (1996)  are shown as solid triangle, and the objects from Patten \\& Pavlovsky (1999) appear as open triangles. Clearly, all these samples merge smoothly, describing a good cluster isochrone. Most of  our probable candidate members are above the empirical  isochrone, which was obtained based on confirmed members of the cluster (Figure 3). In fact, they are located in a  wide band,  slightly larger than $\\sim$0.75 mag (the maximum  shift due to binarity).  We have compared our optical data of IC2391 candidates with data from two other  very well known clusters. Figure 7b contains photometry from the Pleiades (asterisks), alpha Per (plus symbols) and IC2391 (solid and open circles for probable and possible members). These open clusters have lithium ages of 125, 85 and 53 Myr, respectively (Stauffer,  Schultz, \\& Kirkpatrick 1998; Stauffer et al. 1999; Barrado y Navascu\\'es, Stauffer, \\& Patten 1999). The large  width  of the IC2391 main sequence, which includes probable and possible members, can be appreciated. However, this scatter is greatly reduced if we only take into account probable members. Then, the scatter is similar to the value present in the other two clusters. This is consistent with our previous estimate of the contamination for the possible members which are located at the blue side of our empirical IC2391 isochrone. Only further spectroscopy (or, eventually, proper motions) will allow us to verify the membership status of these objects. The comparison between the MS lower envelopes  of these two clusters with the lower MS of IC2391 probable members suggests that, indeed, IC2391 is slightly younger than Alpha Per, and that both of them are considerably younger than the Pleiades."
    },
    "0011/astro-ph0011171_arXiv.txt": {
        "abstract": "s{ We investigate the possibility that ultra high energy cosmic rays ($E>10^{19}$ eV) are related to the distribution of matter on large scales. The large scale structure (LSS) data stems from the recent \\iras\\ \\pscz\\ redshift survey. We present preliminary predictions drawn from   an anisotropic distribution of sources which follows the galaxy distribution. } ",
        "introduction": "\\vspace{-8cm} \\begin{flushright}FISIST/12-2000/CENTRA \\end{flushright} \\vspace{8cm} Ultra high energy cosmic rays (UHECR) are particles with kinetic energies above $\\sim 10^{18}$ eV.\\cite{nagano} The nature of these energetic particles is presently unknown. The reason is twofold: {\\it i)} These particles interact on the top of the Earth's atmosphere producing extensive air showers (EAS) that can be observed from the ground; in this case, the primary particle can not be observed directly; and {\\it ii)} at these ultra high energies (UHE) the fluxes are extremely low (less than 1 particle per square kilometer per year). This makes it impracticable to observe the UHE particles directly using balloons, satellites or spacecrafts due to their small acceptance.  The observation methods are indirect and rely on the observation of the secondary particles produced in the EAS. The hadronic particles, as well as muons and electrons created by the interactions of the primary particle in the atmosphere, are detected on the ground. The EAS also produces detectable fluorescent light photons due to the excitation of nitrogen molecules in the air by the charged secondary particles. Yet, the secondary charged particles that travel with velocities higher than the velocity of the light in the air generate Cherenkov radiation that can also be detected. Despite the fact that the EAS can be observed by the detection of different kinds of secondaries using various techniques, the determination of the nature of the primary particle is very difficult and model dependent. As the fluxes are low it is necessary to use large ground arrays, and/or many fluorescent light and Cherenkov radiation detectors. Until the present moment, the ground arrays and the fluorescent light detectors have gathered only a handful of events in the UHE range. The number of events has not been enough to tell us whether the sources are extragalactic or are located in the Galaxy.  If the primary particle is a $\\gamma$-ray or a neutrino the arrival direction would point back directly to the source. This would also be the case for charged particles if the Galactic magnetic field is $\\lesssim 10^{-6}$ G and the extragalactic magnetic fields are $\\lesssim 10^{-9}$ G in the case of extragalactic sources. The distance to the source is also constrained for most kinds of primaries: If the primary particle is a nucleus or a proton (antiproton) the distance to the source is limited to less than $\\sim 100$ Mpc for particles with arrival energies above $\\sim 6 \\times 10^{19}$ eV (GZK cutoff, see Fig.(1)). Due to interactions with the cosmic microwave background (CMB) photons these particles rapidly lose energy. Sources of $\\gamma$-rays must be even closer because of the short mean absorption length for the UHE photons traveling in the CMB.  The mechanism that provides particles with UHE is also not known. The ignorance about the sources makes it harder to determine the mechanism at work. For the UHE events no source candidate in the vicinity of the region to where the arrival direction points back has been found yet. On the other hand, some analysis of the showers profile have been favoring protons as the primary particle.\\cite{ave} As it was mentioned above, the number of UHE events is still too small to allow us, based on the statistics, to answer questions about their isotropy and composition.  In this work, we  assume that the UHECR primary particles above $10^{19}$ eV are predominantly extragalactic protons and that the sources are related to the distribution of matter on large scales. It means that without specifying the sources themselves or the acceleration mechanism, we would expect an excess of events coming from regions with mass overdensities and less events coming from regions with mass underdensities. In the \\S\\ 2 we briefly describe the formalism used; the propagation code is described in the \\S\\ 3; and the  smoothing procedure of the density field is described in \\S\\ \\ref{density_field}. Our results are presented in the \\S\\ 5. ",
        "conclusions": ""
    },
    "0011/astro-ph0011400_arXiv.txt": {
        "abstract": "Metcalfe et al. (1995, 1996) have shown that galaxy counts from the UV to the near-IR are well-fitted by simple evolutionary models where the space density of galaxies  remains constant with look-back time  while the star-formation rate rises exponentially. We now extend these results, first by using data from the Herschel Deep Field to show that these same models give detailed fits to the faint galaxy $r-i:b-r$ colour-colour diagram. We then use these models to predict the number counts of high redshift galaxies detected by the Lyman break technique. At $z\\approx3$ there is almost exact agreement between our prediction and the data, suggesting that the space density of galaxies at $z\\approx3$ may be close to its local value. At $z\\approx4$ the space density of bright galaxies remains unchanged; however, the space density of dwarf galaxies is significantly lower than it is locally, suggesting that we have detected an epoch of dwarf galaxy formation at $z\\approx4$. Finally, significant numbers of Lyman-break galaxy candidates are also detected at $z\\approx6$ in the Hubble and Herschel Deep Fields; taking this observation together with a number of recent detections of spectroscopically confirmed $z\\approx6$ galaxies suggests that the space density of bright galaxies at  $z\\approx6$ remains comparable to the local space density and thus that the epoch of formation of bright galaxies may lie at yet higher redshift. ",
        "introduction": "The epoch of galaxy formation has been the subject of much recent controversy.  Initially, observational evidence for a late epoch of galaxy formation appeared to come from analysis of the Canada-France Redshift Survey and the Hubble Deep Field where the universal star-formation rate (SFR) appeared to peak at $z\\approx1$ (Lilly et al. 1995, Madau et al. 1996). However, it was also suggested that the CFRS/HDF data could still be consistent with models (Bruzual \\& Charlot 1993) with a constant space density of galaxies whose SFR rises exponentially to redshifts $z\\ga2$ (Metcalfe et al. 1996, Shanks et al. 1998). Indeed, if these `pure luminosity evolution' models are modified by assuming a small amount of dust internal to spirals and  a dwarf-dominated IMF for early-type galaxies, then   they  appeared to give very good fits to galaxy number count data from the UV to the near-infrared, at least in the case of low $\\Omega_0$ (Metcalfe et al. 1996). The recent detection of high redshift galaxies as sub-mm sources (Smail et al. 1997, Hughes et al. 1998, Barger et al. 1998) also appears to support the idea that the SFR may rise to much higher redshifts. The `pure luminosity evolution' models with internal spiral dust absorption have also been shown to be capable of producing fits to the sub-mm galaxy counts and the optical/near-IR galaxy counts simultaneously (Busswell \\& Shanks 2000). Details of further constraints from the UV colours of galaxies on the dust extinction law  and its evolution with redshift are given by Metcalfe et al. (2000).  In the Einstein-de Sitter Universe, with its smaller volume, it has been known for some time that such simple models only fit the optical counts to B$\\la$25mag (Yoshii \\& Takahara 1988, Guiderdoni \\& Rocca-Volmerange 1990, Koo et al. 1993) before underestimating the counts at fainter magnitudes. In this case, a new population of faint galaxies must be postulated; these extra galaxies could be the result of a process of dynamical merging at high redshift (Kauffmann \\& Charlot 1998). Alternatively, as we assume here, they could be due to a population of `disappearing dwarfs', denoted `dE', that are luminous at high redshift but then dim dramatically to be virtually undetectable at low redshift (Babul \\& Rees 1992, Metcalfe et al. 1996, 2000)  \\begin{figure} \\plotone{fig1.eps} \\caption{ Lower points; HST STIS Hubble Deep Field South galaxy counts in the Clear passband from the data of Williams et al. (2000), analysed by techniques described elsewhere (Metcalfe et al. 2000). Also shown are the HST WFPC-2 R(F606W) counts in the HDF-N,-S fields; the STIS counts are slightly lower than, but still statistically consistent with, these data in the range of overlap. Upper points; HST NICMOS HDF-N,-S galaxy count data in the infra-red H (F160W) passband from our analysis (Metcalfe et al. 2000) of the data of Williams et al. (2000) in the case of HDF-S and directly from Thompson et al. (1999) in the case of HDF-N.  The counts have been scaled upwards by a factor of 100 for clarity. HST NICMOS data from other authors are also shown (Teplitz et al. 1998, Yan et al. 1998). The WHDF data comes from a 20hr exposure at the Calar Alto 3.5m telescope using the $\\Omega'$ camera. The field-of-view is 7$'\\times$7$'$ and the galaxy counts reach H=22.5. The models shown are those of Metcalfe et al. (1996) applied to both passbands (see text).} \\end{figure} ",
        "conclusions": "Thus the picture of galaxy formation which is suggested is that bright galaxies form first at $z\\ga4$ and then faint galaxies form later at $3\\la z\\la4$. This is similar to the `downsizing' scenario suggested in other data at $z\\approx1$ (Cowie et al. 1997) and may possibly be connected to claims of evolving, faint luminosity function slopes at even lower redshifts (Brinchmann et al. 1998).  This scenario is unexpected in hierarchical galaxy formation theories where dwarf galaxies are expected to form {\\it earlier} than giant galaxies. In a CDM model with feedback, this order may be reversed with supernovae winds blowing gas out of dwarf galaxies preferentially because of their shallower gravitational potentials. However, suppression of star formation at $z\\ga1$ is usually required in such models to produce the observed luminosity function at the present day (Brinchmann et al. 1998); it thus remains to be seen how well a CDM model will deal with the large numbers of luminous galaxies detected between $4\\la z\\la 7$. New observations of the galaxy luminosity function at even higher redshifts are now required to determine the redshift of formation of the brightest galaxies."
    },
    "0011/astro-ph0011546_arXiv.txt": {
        "abstract": "The LMC, SMC, and the Milky Way are a spectacular set of interacting galaxies. The prominent tidal arms which protrude from the Magellanic Clouds extend for tens of kiloparsecs and cover a large part of the southern sky. These arms, and the Magellanic Clouds themselves, are therefore a highly suitable subject for the detailed investigation of an interacting galaxy system. We present the first results of the Parkes narrow-band \\HI\\ survey of the tidal arms of the Magellanic System. It is the first fully sampled, spatially complete survey of this system.  The data provide information on the column density distribution and the variation of the shape of the line profiles over the extent of the tidal features.  We present a cloud catalogue of the tidal features of the Magellanic System and discuss the distribution relative to the observational parameters longitude, latitude and radial velocity (v$_{\\rm LSR}$ and v$_{\\rm GSR}$).  We find a new stream-like feature that is most likely related to the Leading Arm. Detailed investigation reveals a number of clouds showing a cometary appearance in the position-velocity representation (so-called head-tail structures).  These head-tail structures are interpreted as clouds that are currently interacting with an ambient medium. ",
        "introduction": "Dieter (1965) surveyed the Galactic poles with the Harvard 60-ft antenna and noted several detections with high radial velocities near the southern Galactic pole which have no counterparts towards the northern pole.  This was the first detection in \\HI\\ 21-cm line emission of what we call today the ``Magellanic Stream'', but she did not identify it as a coherent structure.  Further observations of the northern part of the Magellanic Stream, called the ``south pole complex'' at this time, were done by Kuilenburg (1972) and Wannier \\& Wrixon (1972). Mathewson, Cleary, \\& Murray (1974) observed the southern part of the Magellanic Stream using the 18-m Parkes telescope. They discovered the connection between the ``south pole complex'' and the Magellanic Clouds and called it the ``Magellanic Stream''.  Later on, several parts of the Magellanic Stream were observed with higher resolution and sensitivity.  We now know that, next to the Milky Way/Sagittarius dwarf interaction, the LMC, the SMC and the Milky Way form the nearest and certainly one of the most spectacular examples of a galaxy interaction. Individual tidal features that are traced by the \\HI\\ 21 cm line are: (1) the Magellanic Bridge (MB), a high column density gas stream (N$_{\\rm HI} \\geq 1 \\times 10^{21}\\,{\\rm cm^{-2}}$ at 14' angular resolution) connecting the stellar bodies of the LMC and SMC (see Fig. 1); (2) the Magellanic Stream (MS), a coherent structure which starts from the Magellanic Clouds and trails over 100$^{\\circ}$ passing the southern Galactic pole.  The column density in the bright parts of the MS is up to N$_{\\rm HI}\\approx 5 \\times 10^{20}\\,{\\rm cm^{-2}}$ at 14' angular resolution; and (3) the Leading Arm (LA, Putman et al.~1998) that was found to protrude from the Magellanic Bridge and the LMC along several clumpy filaments. The existence of the LA was anticipated by interaction models but not previously detected as a coherent tidal structure.  There are two major approaches to explain these features. The first is a pure gravitational model (e.g.~Gardiner, Sawa, \\& Fujimoto 1994), and the second is a ram-pressure model (e.g.~Moore \\& Davis 1994). The first model is able to explain the existence of a LA feature, while the ram-pressure model fits the morphology of the MS much better, but fails in explaining the LA. Recent models combine both approaches by introducing drag forces to the gravitational model (Gardiner 1999).  Recently, three large scale surveys in the southern sky have been made. The first is the Argentinian survey (\\HI\\ Southern Sky Survey, HISSS), presented by Bajaja et al.~in these proceedings.  It is the counterpart to the Leiden/Dwingeloo survey and offers an angular resolution of 30' with grid spacings of 30'. At a velocity resolution of $\\Delta v$ = 1 km ${\\rm s}^{-1}$ the rms noise is $\\sigma_{\\rm rms} \\approx $ 0.07 K.  The second is the HIPASS survey, also presented in these proceedings. It offers a velocity resolution of $\\Delta v$ = 18 km ${\\rm s}^{-1}$ and an rms noise of $\\sigma_{\\rm rms} \\approx $ 0.01 K. The observing mode of the HIPASS is in-scan beam-switching. This mode filters out all large-scale structure of the Milky Way or the Magellanic System.  The northern part of the Magellanic Stream is not (yet) complete (it is complete for $\\delta_{2000} <$ +2\\deg).  The third survey is presented in this paper. Our survey was performed using the narrow-band facility of the 64-m Parkes telescope. It is fully sampled and covers the complete Magellanic System offering a velocity resolution of $\\Delta v$ = 0.8 km ${\\rm s}^{-1}$ and an rms noise of $\\sigma_{\\rm rms} \\approx $ 0.12 K.  Our narrow-band \\HI\\ survey of the Magellanic System comprises full sampling and high velocity resolution. Both are necessary for a detailed investigation of the physical condition of the gas in the Magellanic System and for improving the constraints which are needed to improve the tidal and ram-pressure interaction models. ",
        "conclusions": "We have presented early results of the first spatially complete and fully sampled \\HI\\ survey of the Magellanic System. We have also compiled a catalogue of tidal gas clouds in the Magellanic System. This catalogue may help to improve tidal models which attempt to reproduce the morphology of the Magellanic System in detail.  We have discovered a new stream-like structure that is most likely a part of the Leading Arm. A future, detailed investigation of the small-scale structure within the tidal arms may address whether there is any interaction with the ambient medium."
    },
    "0011/astro-ph0011293_arXiv.txt": {
        "abstract": "We present a numerical study of topological descriptors of initially Gaussian and scale-free density perturbations evolving via gravitational instability in an expanding universe. We carefully evaluate and avoid numerical contamination in making accurate measurements on simulated fields on a grid in a finite box. Independent of extent of non linearity, the measured Euler number of the excursion set at the percolation threshold, $\\delta_{\\sc c}$, is positive and nearly equal to the number of isolated components, suggesting that these structures are trees.  Our study of critical point counts reconciles the clumpy appearance of the density field at $\\delta_{\\sc c}$ with measured filamentary local curvature. In the Gaussian limit, we measure $|\\delta_{\\sc c}|> \\sigma$ in contrast to widely held belief that $|\\delta_{\\sc c}| \\sim \\sigma$, where $\\sigma^2$ is the variance of the density field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/hep-ph0011073_arXiv.txt": {
        "abstract": "In this talk I discuss models in which a homogeneous scalar field is used to modify standard cosmology above the nucleosynthesis scale to provide an explanation for the observed matter-antimatter asymmetry of the Universe. ",
        "introduction": "Scalar fields are used to model either a very early universe ({\\it inflation}), or a very late universe ({\\it `quintessence'})~\\cite{jf-zws}. Their role at intermediate times is however largely ignored. In this talk~\\footnote{Invited talks presented at the conferences {\\it Strong and Electroweak Matter} (SEWM-2000) in Marseille, (June 14-17, 2000), and {\\it Cosmology and Astroparticle Physics (CAPP-2000)} in Verbier, Switzerland (July 17-28, 2000); based on work with Michael Joyce \\cite{JoyceProkopec,Prokopec,JoyceProkopec2}.} I discuss models in which a homogeneous scalar field is used to modify the standard cosmology at the electroweak scale to allow for baryogenesis, even when the electroweak transition is smooth or weakly first order.  In order to produce any baryon number, a {\\it source} is required that drives the Universe out of equilibrium~\\cite{ShaposhnikovRubakov}. Since at the electroweak scale the expansion rate is very small ($H/T\\sim 10^{-16}$), it is often assumed that it cannot drive baryogenesis, simply because the produced baryon-to-entropy ratio, $n_B/s\\propto H/T$, is too small to account for observation, $(n_B/s)_{\\rm observed}\\sim 5\\times 10^{-11}$. ",
        "conclusions": ""
    },
    "0011/hep-ph0011245_arXiv.txt": {
        "abstract": "We discuss the potential of the Sudbury Neutrino Observatory (SNO) to constraint the four--neutrino mixing schemes favoured by the results of all neutrino oscillations experiments. These schemes allow simultaneous transitions of solar $\\nu_e's$ into active $\\nu_\\mu$'s, $\\nu_\\tau$'s and sterile $\\nu_s$ controlled by the additional parameter $\\cos^2(\\vartheta_{23}) \\cos^2(\\vartheta_{24})$ and they contain as limiting cases the pure $\\nu_e$--active and $\\nu_e$--sterile neutrino oscillations. We first obtain the solutions allowed by the existing data in the framework of the  BP00 standard solar model and quantify the corresponding predictions for the Charged Current  and the Neutral Current to Charged Current event ratios at SNO in the different allowed regions as a function  of the active--sterile admixture. Our results show that some information on the value of $\\cos^2(\\vartheta_{23}) \\cos^2(\\vartheta_{24})$ can be obtained by the first SNO measurement of the CC ratio, while considerable improvement on the knowledge of this mixing will be achievable after the measurement of the NC/CC ratio. ",
        "introduction": "The Sudbury Neutrino Observatory \\cite{sno} is a second generation water Cerenkov detector using 1000 tonnes of heavy water, D$_2$O, as detection medium. SNO was designed to address the problem of the deficit of solar neutrinos observed previously in the Homestake~\\cite{chlorine}, Sage~\\cite{sage}, Gallex+GNO~\\cite{gallex,gno} , Kamiokande~\\cite{kamioka} and Super--Kamiokande~\\cite{superk,suzuki} experiments, by having sensitivity to all flavours of neutrinos and not just to $\\nu_e$, allowing for a model independent test of the oscillation explanation of the observed deficit.  Such sensitivity can be achievable because energetic neutrinos can interact in the D$_2$O of SNO via three different reactions. Electron neutrinos may interact via the Charged Current (CC) reaction \\begin{equation} \\nu_e + d \\to p + p + e^-\\ , \\label{CCreac} \\end{equation} with an energy threshold of several MeV. All non-sterile neutrinos may also interact via Neutral Current (NC) \\begin{equation} \\label{eq:nunc} \\nu_x + d \\to n + p + \\nu'_x  \\ ,\\quad (x=e,\\mu,\\,\\tau), \\label{NCreac} \\end{equation} with an energy threshold of 2.225 MeV. With smaller cross section the non-sterile neutrinos can also interact via Elastic Scattering (ES) $\\nu_x + e^- \\to \\nu'_x + e^- $.  The main objective of SNO is to measure the ratio of NC/CC events. In its first year of operation SNO is concentrating on the measurement of the CC reaction rate while in a following phase, after the addition of MgCl$_2$ salt to enhance the NC signal, it will also perform a precise measurement of the NC rate. It is clear that a cross-section-normalized and acceptance-corrected ratio higher than 1 would strongly indicate the oscillation of $\\nu_e$ into $\\nu_\\mu$ and/or $\\nu_\\tau$. On the other hand a deficit on both CC and NC leading to a normalized NC/CC ratio $1$, can only be made compatible with the oscillation hypothesis if $\\nu_e$ oscillates in to a sterile neutrino.  There are several detailed studies in the literature of the potential of the SNO experiment to discriminate between the different oscillation solutions to the solar neutrino problem (SNP) \\cite{bali,snoothers,bkssno,ournew}. Most of these studies have been performed in the framework of oscillations between two neutrino states where $\\nu_e$ oscillates into either an active, $\\nu_e\\to\\nu_a$,  or a sterile, $\\nu_e\\to\\nu_s$, neutrino channel. On the other hand, once the possibility of a sterile neutrino is considered, these two scenarios are only limiting cases of the most general mixing structure \\cite{DGKK-99,ourfour} which permits simultaneous $\\nu_e\\to\\nu_s$ and $\\nu_e\\to\\nu_a$ oscillations.  In this paper we study the potential of the Sudbury Neutrino Observatory to discriminate between active and sterile solar neutrino oscillations when analyzed in the framework of four--neutrino mixing. We consider those four--neutrino schemes favoured by considering together with the solar neutrino data, the results of the two additional evidences pointing out towards the existence of neutrino masses and mixing: the atmospheric neutrino data \\cite{atmos} and the LSND results \\cite{lsnd}. We concentrate on two SNO measurements: the first expected result on the CC ratio and the expected to be most sensitive, the ratio of NC/CC. The measurement of other observables, such as the recoil energy spectrum of the CC events and the zenith angular dependence \\cite{bali,snoothers,bkssno,ournew} can provide important information to distinguish between the different allowed regions for $\\nu_e$--active oscillations but they are not expected to be very sensitive as discriminatory between the active and sterile oscillations.  The outline of the paper is the following. For the sake of completeness we begin by discussing in  Sec.~\\ref{two} the expected results when obtained in the pure two--neutrino oscillation hypothesis. In Section~\\ref{fourana} we determine the presently allowed regions for the oscillation solutions to the SNP  in the framework of four--neutrino mixing. In Sec.~\\ref{foursno} we present the results of the expected CC and NC/CC rates for the different solutions and quantify the attainable sensitivity to the additional mixing controlling the admixture of active--sterile entering into the solar neutrino oscillations. Finally in Sec.~\\ref{discuss} we summarize our conclusions. ",
        "conclusions": "\\label{discuss}  In this paper we have studied the potential of the Sudbury Neutrino Observatory to discriminate between active or sterile solar neutrino oscillations when analyzed in the framework of four--neutrino mixing. We considered those four--neutrino schemes favoured by considering together with the solar neutrino data, the results of the two additional evidences pointing out towards the existence of neutrino masses and mixing: the atmospheric neutrino data \\cite{atmos} and the LSND results \\cite{lsnd}. These schemes allow simultaneous transitions of solar $\\nu_e's$ into active $\\nu_\\mu$'s, $\\nu_\\tau$'s and sterile $\\nu_s$ controlled by the additional parameter $\\cos^2(\\vartheta_{23}) \\cos^2(\\vartheta_{24})$ and they contain as limiting cases the pure $\\nu_e$--active and $\\nu_e$--sterile neutrino oscillations. The allowed solar solutions have been reanalyzed including the recently BP00 standard solar model and the latest solar neutrino data. We find that the global minimum lies in the LMA region and for pure $\\nu_e$--active oscillations ($c_{23}^2c_{24}^2=0$). We also find that in the framework of four--neutrino mixing the  SMA solution is allowed at 90\\% CL for non vanishing active--sterile mixing $c_{23}^2c_{24}^2$ in the range [0.11,0.31].  We concentrated on two SNO measurements: the first expected result on the [CC] ratio and the expected to be most sensitive to the active--sterile admixture, the ratio of [NC]/[CC] and evaluated the predictions in the different regions as a function of the additional mixing $c_{23}^2c_{24}^2$. Our results are display in Figs.~\\ref{4fsno_sma}$-$\\ref{4fsno_low}. They show that in most cases with the measurement of the [CC] ratio, it will be hard to improve the present knowledge of $c_{23}^2c_{24}^2$ but with the precise determination of the [NC/CC] ratio at SNO, this parameter can be strongly constrained for some of the allowed solutions. For example, we find that for the [CC] rate the most discriminatory scenario would be that SNO finds a small value [CC]$\\sim 0.25$. This significantly hints towards the LMA solution to the solar neutrino problem and towards an active--active  oscillation scenario. In this case the [NC]/[CC]$\\sim$ 4 ($\\pm$ 0.7 at $5\\sigma$) will be conclusive for selecting LMA as the solution to the SNP and will imply and upper bound on  $c_{23}^2c_{24}^2<0.3$. Conversely, a measurement of  [NC]/[CC]$\\sim 1$, although harder to distinguish from the non--oscillation scenario, will push the oscillation hypothesis towards the sterile SMA solution, and with the expected sensitivity a bound $c_{23}^2c_{24}^2>0.44$ at 99\\% CL can be imposed."
    },
    "0011/astro-ph0011016_arXiv.txt": {
        "abstract": "We have collected a large body of NIR ($H$ band), UV (2000 \\AA) ~and H$\\alpha$ measurements of late-type galaxies. These are used, jointly with spectral evolutionary synthesis models, to study the initial mass function (IMF) in the mass range $m$ $>$ 2 ${\\rm M\\odot}$. For spirals (Sa-Sd), Magellanic irregulars (Im) and blue compact dwarfs (BCD), our determination is consistent with a Salpeter IMF with an upper mass cutoff $M_{up}$ $\\sim$ 80 ${\\rm M\\odot}$. \\noindent The history of star formation and the amount of total gas (per unit mass) of galaxies are found to depend primarily on their total masses (as traced by the $H$ band luminosities) and only secondarily on morphological type. The present star formation activity of massive spirals is up to 100 times smaller than that average over their lifetime, while in low mass galaxies it is comparable to or higher than that at earlier epochs. Dwarf galaxies have presently larger gas reservoirs per unit mass than massive spirals. \\noindent The efficiency in transforming gas into stars and the time scale for gas depletion ($\\sim$ 10 Gyrs) are independent of the luminosity and/or of the morphological type. \\noindent These evidences are consistent with the idea that galaxies are coeval systems, that they evolved as closed-boxes forming stars following a simple, universal star formation law whose characteristic time scale is small ($\\tau$ $\\sim$ 1 Gyr) in massive spirals and large ($\\tau$ $>$ 10 Gyr) in low mass galaxies. A  similar conclusion was drawn by Gavazzi \\& Scodeggio (1996) to explain the colour-magnitude relation of late-type galaxies. \\noindent The consequences of this interpretation on the evolution of the star formation rate and of the gas density per comoving volume of the Universe with look-back time are discussed. ",
        "introduction": "The hierarchical (Cole et al. 1994) and the monolithic collapse (Sandage 1986) scenarios of galaxy formation make different predictions on the evolution of the primeval density perturbations which gave birth to galaxies. While in hierarchical models small objects form first and large objects are formed via successive merging of smaller structures, in the monolithic scenario galaxies are formed via a unique collapse of the original density perturbation without any interaction with the environment (closed box). A wide variety of physical and morphological properties observed in nearby galaxies are expected from both models: while in a monolithic scenario ellipticals should have formed most of their stars in a time short compared to the collapse time in order to avoid the formation of a disc by gas-gas dissipation, in spirals the slower initial star formation activity allowed a significant fraction of gas to form a rotating disc (Sandage 1986). In hierarchical models the different properties of galaxies are the result of the initial conditions of the original merger.  The predicted luminosity functions and density distributions of galaxies in the Universe should in principle allow to disentangle observationally between these two models (Kauffmann et al. 1999). Furthermore these models make predictions on the time evolution of the physical properties of galaxies, such as their structural parameters (e.g. bulge-to-disk ratio), their dynamics, their star formation history (thus their stellar populations), their gaseous content and metal enrichement (Kauffmann \\& Charlot 1998 and references therein). An accurate determination of the dependence of these properties on $z$ (e.g. Cowie et al. 1996) will allow us to discriminate between the galaxy formation scenarios.\\\\ A detailed knowledge of the phenomenology of galaxies at $z$=0, which represents the present stage of galaxy evolution, is of primary importance for constraining evolutionary models.\\\\ Ironically the \"zero point\" of galaxy evolution is still not sufficiently well determined, due to the lack of extensive observations of local galaxies. For example the systematics of their current star formation activity, gaseous content and metallicity are still poorly known. In particular the UV and H$\\alpha$ observations which are representative of the current star formation activity of galaxies (Donas et al. 1987, Kennicutt 1998), the near-IR data which trace the old stellar population, and spectro-photometric measurements (Zaritsky et al. 1994) which are necessary to derive the properties of the underlying stellar population and the physical conditions of the interstellar medium, are still rare.\\\\ With the ambitious aim of constructing a representative description of the physical properties of nearby galaxies, we undertook a multi-frequency survey of $\\sim$ 3500 optically selected galaxies spanning the broadest possible range in morphological type (Ellipticals, Spirals, dE, Im and BCD), luminosity (-22 $\\leq$ $M_B$ $\\leq$ -13), and environmental conditions (clusters - isolated). Our own H$\\alpha$, near-IR, millimetric and centimetric observations, together with data from the literature, were collected in a multifrequency database (see Gavazzi, Boselli \\& Pierini, 1996).\\\\ This database has been used to interpret the colour-magnitude relation (Gavazzi \\& Scodeggio 1996) and  the star formation activity of late-type galaxies (Gavazzi et al. 1998). The extension of our near-IR survey to faint spirals, Im and and blue compact dwarfs (BCD), which is the subject of the present series of papers, combined with recent H$\\alpha$ observations of a set of late-type galaxies, and with the large body of available HI and CO radio data makes it possible to study, in the present work, the phenomenology of the young stellar population and of the gas content of a large sample of late type galaxies from Sa to Im-BCD.\\\\   \\noindent Previous works devoted at analysing the phenomenology of nearby galaxies have shown that the fraction of available gas and the activity of star formation (normalized to the visible stellar luminosity) increases along the Hubble sequence (Roberts \\& Haynes 1994; Kennicutt 1998). However none of these works was focused on the role of the total mass in governing the evolution of galaxies, an issue that we address in the present work. We do so using the observation that the near-IR $H$ luminosity, which is dominated by the emission of the old main sequence and red giant stars (Bruzual \\& Charlot 1993), is a direct tracer of the dynamical mass of late-type systems (within the optical radius; Gavazzi et al. 1996a). Since the near-IR mass to light ratio was found not to vary with mass for galaxies spanning the whole range in morphological type from Sa to Im and BCD, contrary to what was found in the optical, by normalizing other observed quantities (e.g. the gas content etc.) to the near-IR luminosity we better remove than using optical luminosities the well known observed luminosity-luminosity or luminosity-mass scaling relations. Furthermore the decomposition of the near-IR light profiles are used to discriminate between disc dominated and bulge dominated galaxies. This new approach is made possible by the availability of a large body of near-IR observations of nearby galaxies described in papers I, II, III and IV of this series. Furthermore the available HI and CO data allow a precise determination of the total gas content for a sample of galaxies with a broad coverage in morphological type and luminosity.  The sample used and the data coverage is described in Section 2. Early-type (type $\\le$ S0a) galaxies are excluded from the present analysis since their UV emission, being dominated by the extreme horizontal branch stars and their progenitors (Dorman et al. 1995), does not give information on the young stellar population. Moreover we restrict the present analysis to galaxies which do not show signs of interaction with their surroundings. For this we use the HI-deficiency parameter to discriminate between \"normal\" galaxies and galaxies suffering for gas depletion due to ram pressure and we include in the present study only isolated objects or cluster galaxies with a HI-deficiency (defined as the ratio of the HI mass to the average HI mass of isolated objects of similar morphological type and linear size (Haynes \\& Giovanelli 1984)) $\\leq 0.3$.\\\\ The current star formation rate (SFR) and the initial mass function (IMF) of late-type galaxies are studied in Section 3. The dependence of the total (atomic plus molecular) gas content as a function of morphological type and luminosity is analysed in Section 4. A model for the star formation history of late-type galaxies is presented in Section 5. The dependence of the SFR and of the gas density of the Universe on look-back time is briefly discussed in Section 6. An appendix describes how the H$\\alpha$ and UV (2000 \\AA) data are corrected for [NII] contamination and internal extinction.  \\section {The sample}  Galaxies analysed in this work are taken from the Zwicky catalogue (CGCG, Zwicky et al. 1961-1968)($m_{pg}$ $\\leq$ 15.7). They are either late-type (type$>$S0a) members of 3 nearby ($vel$ $\\le$ 8000 km s$^{-1}$) clusters (Cancer, A1367, Coma), or located in the relatively low-density regions of the Coma-A1367 supercluster (11$^h$30$^m$ $<$RA$<$ 13$^h$30$^m$; 18$^o$$<$dec$<$32$^o$) as defined in Gavazzi et al. (1999a). To extend the present study to lower luminosities, we include in the sample the late-type Virgo cluster galaxies brighter than $m_{pg}$ $\\leq$ 14.0 listed in the Virgo Cluster Catalogue as cluster members (VCC, Binggeli et al. 1985). Furthermore VCC galaxies with 14.0 $\\leq$ $m_{pg}$ $\\leq$ 16.0 included in the \"ISO\" subsample described in Boselli et al. (1997a) and CGCG galaxies in the region 12$^h$ $<$RA$<$ 13$^h$; 0$^o$$<$dec$<$18$^o$ but outside the VCC, are considered. To avoid systematic environmental effects we consider the subsample of late-type galaxies whose HI-deficiency is $\\leq$0.3, typic of unperturbed, isolated galaxies. The final combined sample comprises 233, mainly ``normal'' galaxies (a few starburst or active galaxies might however be included). \\\\ The accuracy of the morphological classification is excellent for the Virgo galaxies (Binggeli et al. 1985; 1993). Because of the higher distance, the morphology of galaxies belonging to the other surveyed regions suffers from an uncertainty of about 1.5 Hubble type bins.\\\\ We assume a distance of 17 Mpc for the members (and possible members) of Virgo cluster A, 22 Mpc for Virgo cluster B, 32 Mpc for objects in the M and W clouds (see Gavazzi et al. 1999b). Members of the Cancer, Coma and A1367 clusters are assumed at distances of 62.6, 86.6 and 92 Mpc respectively. Isolated galaxies in the Coma supercluster are assumed at their redshift distance adopting $H_o$ = 75 km s$^{-1}$ Mpc$^{-1}$.\\\\ For the 233 optically selected galaxies, complementary data are available in other bands as follows: 100\\% have HI (1420 MHz), 99\\% $H$ band (1.65 $\\mu$m) data available. A much sparser coverage exists in the UV (2000 \\AA)(29 \\%), CO (115 GHz)(38\\%) and H$\\alpha$ (6563 \\AA) (65\\%), as shown in Table 1 and 2.  The morphological distribution of galaxies is given in Table 3.  \\subsection {Data analysis}  H$\\alpha$+[NII] fluxes obtained from imaging, aperture photometry or integrated spectra are taken from Kennicutt \\& Kent (1983), Kennicutt et al. (1984), Gavazzi et al. (1991) Gavazzi et al. (1998), Young et al. (1996) Almoznino \\& Brosch (1998), Moss et al. (1998), Heller et al. (1999) and references therein. H$\\alpha$ fluxes from Kennicutt \\& Kent (1983), Kennicutt et al. (1984) and Gavazzi et al. (1991) have been multiplied by 1.16, as suggested by Kennicutt et al. (1994), in order to account for the continuum flux overestimate due to inclusion of the telluric absorption band near 6900 \\AA ~in the comparison filter. Additional observations of 66 galaxies have been recently obtained by us during several runs at the Observatoire de Haute Provence (France), at San Pedro Martir (Mexico) and at Calar Alto (Boselli et al. in preparation; Gavazzi et al. in preparation). The estimated error on the H$\\alpha$+[NII] flux is $\\sim$ 15\\%.  The UV data (at 2000 \\AA) are taken from the FOCA (Milliard et al. 1991) and FAUST (Lampton et al. 1990) experiments; UV magnitudes are from Deharveng et al. (1994) and Donas et al. (1987; 1990; 1995; in preparation). The estimated error on the UV magnitude is 0.3 mag in general, but it ranges from 0.2 mag for bright galaxies to 0.5 mag for weak sources observed in frames with larger than average calibration uncertainties.  HI data are taken from Scodeggio \\& Gavazzi (1993) and Hoffman et al. (1996, and references therein). HI fluxes are transformed into neutral hydrogen masses with an uncertainty of $\\sim$ 10\\%.  CO data, used to estimate the molecular hydrogen content, are from Boselli et al. (1997b), Boselli et al. (1995) and references therein. The average error on CO fluxes is $\\sim$ 20\\%; the error on the H$_2$ content, however, is significantly larger (and difficult to quantify) due to the poorly known CO to H$_2$ conversion factor (see Boselli et al. 1997b).  NIR data, mostly from Nicmos3 observations, are taken from this series of papers: Gavazzi et al. (1996b,c: Paper I and II), Boselli et al. (1997a), Gavazzi et al. (2000a, Paper III), Boselli et al. (2000, Paper IV). From these data we derive total (extrapolated to infinity) magnitudes $H_T$ determined as described in Gavazzi et al. (2000b: Paper V) with typical uncertainties of $\\sim$ 10 \\%. These are converted into total luminosities using: $log L_H = 11.36 - 0.4H_T +2logD$ (in solar units), where $D$ is the distance to the source (in Mpc). For a few objects we derive the $H$ luminosity from $K$ band measurements assuming an average $H-K$ colour of 0.25 mag (independent of type; see Paper III).  A minority of the objects in our sample have an $H$ band magnitude obtained from aperture photometry, thus with no asymptotic extrapolation. For these we use the magnitude $H_{25}$ determined as in Gavazzi \\& Boselli (1996) at the optical radius (the radius at which the $B$ surface brightness is 25 mag arcsec$^{-2}$) which is on average 0.1 magnitudes fainter than $H_T$ (Gavazzi et al. 2000a,b).  The total $H$ magnitudes are corrected for internal extinction according to Gavazzi \\& Boselli (1996). No correction has been applied to galaxies of type $>$ Scd. The model independent near-IR concentration index parameter $C_{31}$, defined as the ratio of the radii containing 75 \\% to 25 \\% of the total light of a galaxy, will be used throughout the paper to discriminate between disc dominated and bulge dominated galaxies. As shown in paper V, pure exponential discs are characterized by $C_{31}$ $<$ 3, while $C_{31}$ is $>$ 3 in galaxies with prominent bulges.  \\section {The SFR in galaxies}  Various indicators of the star formation activity of late-type galaxies have been proposed in the literature (see for a review Kennicutt 1998). In the present analysis we use the H$\\alpha$ and UV (2000 \\AA) luminosities, which are commonly accepted as the most direct indicators of the star formation rate (SFR).  \\subsection {SFR from H$\\alpha$ and UV (2000 \\AA) luminosities}  The H$\\alpha$ luminosity gives a direct measure of the global photoionization rate of the interstellar medium due to high mass ($m$ $>$ 10 ${\\rm M\\odot}$), young ($\\le$ 10$^7$ years) O-B stars (Kennicutt 1983; 1990; Kennicutt et al. 1994). The total SFR can be determined by extrapolating the high-mass SFR to lower mass stars using an assumed initial mass function (IMF) $\\psi(m)$:  \\begin{equation} {\\psi (m)=\\int^{M_{up}}_{M_{low}} k~m^{-\\alpha}dm } \\end {equation}  \\noindent where $M_{up}$ ($M_{low}$) is the upper (lower) mass cutoff and $\\alpha$ the slope of the IMF. In the assumption that the SFR is constant on a time scale of some 10$^7$ years, the SFR (in ${\\rm M\\odot yr^{-1}}$) is given by the relation:  \\begin{equation} {SFR_{H\\alpha} ({\\rm M\\odot yr^{-1}}) = K_{H\\alpha}(\\alpha,M_{up},M_{low}) L_{H\\alpha} ({\\rm erg~s^{-1}})} \\end {equation}  \\noindent where $K_{H\\alpha}(\\alpha,M_{up},M_{low})$ is the proportionality constant between the H$\\alpha$ luminosity $L_{H\\alpha}$ (in erg s$^{-1}$) and the $SFR_{H\\alpha}$ (in ${\\rm M\\odot yr^{-1}}$). The value of $K_{H\\alpha}(\\alpha,M_{up},M_{low})$, which depends on the slope $\\alpha$ and on the upper and lower mass cutoffs $M_{up}$ and $M_{low}$ of the IMF, can be determined from models of stellar population synthesis. Different values of $K_{H\\alpha}(\\alpha,M_{up})$ from the stellar population synthesis models of Charlot \\& Fall (1993) are given in Table 4, all for a lower mass cutoff $M_{low}$ = 0.1 ${\\rm M\\odot}$ and a solar metallicity.   The UV emission of a galaxy at 2000 \\AA~ is dominated by the emission of less recent ($\\sim$ 10$^8$ years) and massive (2$<$ $m$ $<$ 5 ${\\rm M\\odot}$) A stars (Lequeux 1988). The UV emission becomes stationary if the SFR is constant over an interval of time as long as the life time of the emitting stars on the main sequence, i.e. $\\geq$ 3 10$^8$ years. Thus, assuming a star formation rate constant on time scales of some 10$^8$ years, the rate of star formation $SFR_{UV}$ from UV luminosities at 2000 \\AA~ ($L_{UV}$) can be determined from the relation:  \\begin{equation} {SFR_{UV} ({\\rm M\\odot yr^{-1}}) = K_{UV}(\\alpha,M_{up},M_{low}) L_{UV} ({\\rm erg~s^{-1}A^{-1}})} \\end {equation}  \\noindent where $K_{UV}(\\alpha,M_{up},M_{low})$ is the proportionality constant between the UV luminosity $L_{UV}$ (in erg s$^{-1}$\\AA$^{-1}$) and the $SFR_{UV}$ (in ${\\rm M\\odot yr^{-1}}$). Values of $K_{UV}(\\alpha,M_{up})$ at 2000 \\AA ~ for $M_{low}$ = 0.1 ${\\rm M\\odot}$ and a solar metallicity kindly made available by S. Charlot are listed in Table 4.  H$\\alpha$ and UV fluxes can thus be used independently to estimate the rate of star formation in galaxies once they are corrected for internal extinction and for the contribution of the [NII] emission line ($\\lambda$ 6548 \\AA~ and 6584 \\AA) which contaminates the narrow band H$\\alpha$ photometry. The corrections applied to the present data are described in the Appendix. As discussed in the Appendix, the lack of far-IR data and of integrated spectra for all the sample galaxies prevents us to make accurate [NII] contamination and extinction correction for each single galaxy, forcing us to use ``statistical'' corrections. Given the broad distribution in the UV and H$\\alpha$ extinctions measured by Buat \\& Xu (1996) and given in Fig. 11 for each morphological class, we estimate an uncertainty in the extinction correction of a factor of $\\sim$ three for each single galaxy, but probably lower for an entire class of objects. This uncertainty, which might be responsible for part of the observed scatter in several star formation indicators, such as the birthrate parameter, should not be critical in the forecoming analysis since it is significantely smaller than the whole range in star formation observed in the sample galaxies, which spans more than 3 orders of magnitude.\\\\ From now on we estimate $SFRs$ (in ${\\rm M\\odot yr^{-1}}$) from UV and H$\\alpha$ luminosities using $K_{UV}(\\alpha,M_{up},M_{low})$ and $K_{H\\alpha}(\\alpha,M_{up},M_{low})$ from Table 4 for $\\alpha$=2.5, $M_{up}$=80 ${\\rm M\\odot}$ and $M_{low}$=0.1 ${\\rm M\\odot}$; for galaxies with both UV and H$\\alpha$ measurements (52 objects), the adopted $SFR$ is an average of the results of the two methods.  \\subsection {The IMF}  Since the H$\\alpha$ emission is due to massive ($m$ $\\geq$ 10 ${\\rm M\\odot}$) O-B stars, and the UV emission to moderate mass (2 $\\leq$ $m$ $\\leq$ 5 ${\\rm M\\odot}$) A stars, the two SFR determinations provide us with an indirect method for studying the IMF of spiral galaxies in the mass range $m$ $\\geq$ 2 ${\\rm M\\odot}$ (Buat et al. 1987; Bell \\& Kennicutt 2000). This method can be applied only if the SFR has been constant over the last 3 10$^8$ years, corresponding to the life time of stars dominating the UV emission on the main sequence. It is well known that this assumption applies for \"normal\" unperturbed objects, such as the ones selected here, while it does not hold for interacting systems (Kennicutt et al. 1987) and starburst galaxies.   For all galaxies in our sample with available H$\\alpha$ and UV data (52 objects), we compare in Fig. 1 the ratio of the H$\\alpha$ to UV fluxes, corrected for extinction and [NII] contamination as described in the appendix, with the values determined from the stellar population synthesis models given in Table 4. Different symbols are use for disc dominated (filled circles) and bulge dominated (empty circles) galaxies. The logarithm of the H$\\alpha$ to UV flux ratio is plotted versus the morphological type and the $H$ luminosity.  The average value $log [flux H\\alpha/flux UV(2000\\AA)]$=1.43 $\\pm$0.25 is consistent with a IMF with slope $\\alpha$ = 2.5 and $M_{up}$ = 80 ${\\rm M\\odot}$, and differs significantly from the values obtained for an IMF with slopes 1.5 or 3.5 (see Table 4). Figure 1 a and b show that $log [flux H\\alpha/flux UV(2000\\AA)]$ is independent of the morphological type, mass and of the presence of a bulge. The value $\\alpha$ = 2.5 is consistent with the Salpeter IMF ($\\alpha$ = 2.35; Salpeter 1955). As shown in Fig. 1, the H$\\alpha$ to UV flux ratio is more sensitive to the slope of the IMF than to $M_{up}$.    As discussed by Calzetti (1999), the H$\\alpha$ to UV flux ratio should increase with extinction in any fixed geometry. In the case of a slab model (absorbing dust and emitting stars well mixed in a disc), the observed $H\\alpha/UV(2800 \\AA)$ flux ratio overestimates the real value by a factor of 2 for typical H$\\alpha$ extinctions of $\\sim$ 1.1 mag (Calzetti 1999). This value has however to be taken as un upper limit since it has been observed that, despite a factor of $\\sim$ 4 in the galactic extinction law between 6563 \\AA ~and 2000 \\AA, UV and H$\\alpha$ extinctions are comparable because of a less efficient absorption of the 2000 \\AA ~ photons dictated by the less concentrated distribution of the UV emitting stars relative to the ionizing stars inside the dusty HII regions (see Appendix). We expect that variations in the extinction are responsible for an important fraction of the scatter in the H$\\alpha$ to UV flux ratio observed in Fig. 1 without introducing any important second order systematic effect (see Appendix).  An $\\alpha$=2.5 is in agreement with the value obtained by Buat et al. (1987), who analysed 31 galaxies with H$\\alpha$ and UV data, adopting a method similar to the one used in this work, coupled with an older version of the evolutionary synthesis models. We do not confirm the trend they find with the morphological type, which might be contributed to by HI-deficient cluster galaxies which we deliberately excluded from our analysis.  \\section {The star formation history of galaxies}  \\subsection{The birthrate parameter b}  The birthrate parameter $b$ is defined by Kennicutt et al. (1994) as the ratio of the current SFR to the average past SFR. This distance independent parameter $b$ is given by:  \\begin{equation} {b=\\frac{SFR}{<SFR>_{past}}=\\frac{SFR ~t_o ~(1-R)}{M_{star}}} \\end{equation}  \\noindent where $t_o$ and $M_{star}$ are the age and the total stellar mass of the disc and $R$ is the fraction of gas that stars re-injected through stellar winds into the interstellar medium during their lifetime. While the SFR comes directly from the H$\\alpha$ and the UV luminosities, the remaining parameters must be estimated indirectly. Since stars eject different fractions of gas into the ISM at different epochs of their life, the return parameter $R$ is a function of time and depends on the assumed IMF and birthrate history. Kennicutt et al. (1994) have shown that 90\\% of the returned gas is released in the first Gyr (more than half in the first 200 Myr) of a stellar generation for any assumed IMF. For this reason we can safely assume an instant recycling gas approximation in the determination of the birthrate parameter $b$, with a constant value $R$=0.3, as determined by Kennicutt et al. (1994) for a Salpeter IMF.  The total stellar mass $M_{star}$ can be determined using the same procedure described in Kennicutt et al. (1994), but using $H$ band luminosities instead of optical ones. Near-IR luminosities are directly proportional to the total dynamical masses ($M_{tot}/L_H$=4.6) at the $B$ band 25 mag arcsec$^{-2}$ isophotal radius, independent of morphological type (Gavazzi et al. 1996a). The same $M_{tot}$ vs. $L_H$ relationship is followed by BCDs, where however the near-IR flux might be contaminated by the emission of red supergiants; thus even in these objects the $H$ luminosity can be properly used as an estimator of the total mass.  Assuming $t_o$ $\\sim$12 Gyrs, the birthrate parameter $b$ comes directly from:  \\begin{equation} {b=\\frac{SFR ~t_o(1-R)}{L_H(M_{tot}/L_H)DM_{cont}}} \\end{equation}  \\noindent where $DM_{cont}$ is the dark matter contribution to the $M_{tot}/L_H$ ratio at the optical radius, that we assume to be $DM_{cont}$=0.5, as in Kennicutt et al. (1994). As discussed by Rubin (1987), this value is independent on type and luminosity and thus it should not introduce any systematic effect in the determination of $b$. Sandage (1986) and Kennicutt et al. (1994) have shown that the birthrate parameter, when determined using $B$ luminosities, increases along the Hubble sequence, but it is still not known how $b$ scales with mass.  The birthrate parameter (in logarithmic units) is plotted in Fig. 2 a and b versus the morphological type and the $H$ luminosity. Fig. 2 shows a strong relationship between the birthrate parameter and these structural parameters, with late-type and dwarf galaxies (low mass objects) all characterized by a similar present and past star formation rate ($b$ $\\sim$ 1). There are a few galaxies with $b$ $>$ 1, i.e. with present star formation rate higher than averaged over the past. These are low-mass, dwarf galaxies which might undergo episodes of star bursts. Massive spirals, on the contrary, have present SFRs significantly lower than in the past ($b$ $\\sim$ 0.1-0.01). The relationship between $b$ and the $H$ luminosity is however considerably clearer than with the Hubble type, whilst only a systematic difference is observed between types $\\leq$ Sbc ($b$ $\\leq$ 0.25) and types $\\geq$ Scd ($b$ $\\geq$ 0.25), with Sc spanning the whole range in $b$.  If we limit our analysis to Virgo cluster galaxies, for which the morphological classification is more accurate, we observe a less dispersed relation between $b$ and the morphological type. A relationship between $b$ and the $H$ luminosity is however still observed inside any given morphological class, implying that part of the observed scatter in the $b$ vs. type relationship is due to a stronger dependence of $b$ on the $H$ luminosity.  On the other hand part of the scatter in the $b$ vs. $log(L_H$) relationship is due to the presence of a bulge; at any luminosity, galaxies with strong bulges ($C_{31}$$>$3; open circles) have lower birthrate parameters than pure exponential discs ($C_{31}$$<$3; filled circles).  \\subsection {The gas content of galaxies}  It is well known that the fraction of atomic gas increases along the Hubble sequence (Roberts \\& Haynes 1994, and references therein), but it is still unknown how the total gas mass (atomic plus molecular) scales with the total mass of galaxies. The line emission of HI at 21 cm and of $^{12}$CO(1-0) at 2.6 mm can be used to estimate the total gas content (HI + H$_2$) of the target galaxies.  The molecular hydrogen content can be estimated assuming a constant ratio between the $^{12}$CO(1-0) line emission and the H$_2$ surface density. In this work we follow Boselli et al. (1997b), but adopting the CO to H$_2$ conversion factor $X$=1.0 10$^{20}$ mol cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ of Digel et al. (1996). For galaxies with no CO measurement, we assume that the molecular hydrogen content is 10\\% of the HI, as estimated from isolated spiral galaxies by Boselli et al. (1997b). The total gas mass is increased by a 30\\% to take into account the helium contribution. The total gas content (normalized to the mass) of a galaxy depends strongly on the morphological type and on the $H$ luminosity as shown in Fig. 3. Late-type, low mass galaxies have a larger amount of gas (per unit mass) than early-type, massive discs.  The relationship between $log(M_{gas}/L_H)$ and Hubble type has a slightly smaller dispersion if limited to the Virgo cluster galaxies which have more reliable morphological classifications. The relationship with the total mass is clearer than with the Hubble type. As for the birthrate parameter, we observe a trend between $log(M_{gas}/L_H)$ and $L_H$ inside each morphological class implying that part of  the scatter in the gas vs. morphology relationship is due to the large scatter in the galactic mass within a given type. The residual of the $log(M_{gas}/L_H)$ vs. $L_H$ relationship, however, is not correlated to the presence of a bulge.  \\subsection {The relation between gas content and the history of star formation in galaxies}  A relationship between the SFR and the gas density is known to exist locally within the discs of nearby late-type galaxies (Kennicutt 1989). Here we show that this relation extends to global quantities, integrated over the whole galaxy. This is not obvious a priori, since it is well known that most of the HI gas reservoir is located outside the optical disc of spiral galaxies, where the star formation does not take place.  Figure 4 shows the relationship between the history of star formation as traced by the birthrate parameter and the total gas reservoir (per unit galaxy mass).  Galaxies with $b$ $\\sim$ 1 have large amount of gas, while objects which formed most of their stars in the past ($b$ $\\sim$ 0.1-0.01) almost exhausted their gas reservoirs, becoming quiescent. Figure 4 shows a segregation between pure discs (filled circles) and bulge galaxies (open circles), the former being star-forming and gas-rich, the latter more quiescent and gas-poor.  \\subsection {The SFE and the gas consumption timescale of galaxies}  The comparison of the current SFR to the gas content (atomic and molecular) gives an estimate of the efficiency of a given object to transform its gas reservoir into stars. The star formation efficiency (SFE) is defined as:  \\begin{equation} {SFE = \\frac{SFR}{M_{gas}} ~~~({\\rm yr^{-1}})} \\end{equation}  \\noindent No relationship is observed between the SFE and the  morphological type or the $H$ luminosity (Fig 5a and b). Pure exponential discs (filled circles) have, on average, higher SFEs than bulge dominated spirals (empty circles) of similar $H$ luminosity. If restricted to disc dominated galaxies, a trend is observed, with low mass galaxies having a lower SFE than giant discs.   The SFE can vary within a given morphological class by up to a factor 10. The average value is 4 10$^{-10}$ yrs$^{-1}$. The SFE is a time independent quantity: it measures the $\\it{instantaneous}$ efficiency of transforming the gas into stars. Since the atomic gas has to go through the molecular phase to form stars inside molecular clouds, a more accurate measure of the SFE can be obtained if $M_{gas}$ in eq. (6) is replaced by the gas $\\it{instantaneously}$ contributing to the process of star formation, thus by $M(H_2)$. Even using this modified definition of the SFE, we do not observe any relationship between the SFE and other structural parameters, in agreement with Rownd \\& Young (1999).  The SFE is proportional to the time-scale for gas depletion if the fraction of gas ejected by stars and recycled is taken into account.  This time-scale, generally referred to as the \"Roberts' time\" (Roberts 1963) is given by the relation (Kennicutt et al. 1994):  \\begin{equation} {\\tau _R = \\frac{(\\frac{M_{gas}}{SFR})}{(1-R)} = \\frac{(\\frac{1}{SFE})}{(1-R)}} \\end{equation}  \\noindent where $R$ is the returned gas fraction. As discussed previously, the $R$ parameter changes with time, with the IMF and with the birthrate history. The determination of the future evolution of the disc is strongly related to the dependence of $R$ on the birthrate history of a galaxy. Kennicutt et al. (1994) have studied how $(1-R)^{-1}$, the correction parameter for the determination of the Roberts' time, changes for different IMFs, star formation laws (i.e. the relationship between the gas surface density and the SFR, known as the Schmidt law) and star formation efficiencies, these three variables being the most important parameters for the determination of the birthrate history of a galaxy. Their analysis have shown that $(1-R)^{-1}$ is in the range 1.5 - 4 for most star forming discs characterized by a birthrate parameter $b$=1-0.1, while it can be larger for rapidly evolving galaxies ($b$ $<$ 0.1) with low gas surface densities. Since the determination of $R$ as a function of the star formation history of a galaxy is behind the scope of the present paper, we assume a standard correction factor $(1-R)^{-1}$ = 2.5, consistent with Kennicutt et al. (1994).  The Roberts' time is independent of the mass and of the morphological type (Fig. 6). This result is in contrast with that obtained by Donas et al. (1987), who observed a weak trend between the Roberts' time and the morphological type, with late-type systems showing a longer $\\tau _R$. The time-scale for gas depletion is in the range between $\\sim$ 1 Gyr and 25 Gyr, with an average value of 10 Gyr, consistent with the value found by Kennicutt et al. (1994). These values are upper limits to the real gas consumption time scales because a large part of the available HI gas is located outside the optical disc of galaxies and thus does not contribute directly to the star formation, unless gas inflow toward the centre of the galaxies is invoked.  \\section {A model of the star formation history}  Let's summarize the evidences collected so far:  \\noindent 1) photometric H$\\alpha$ and UV (2000 \\AA) measurements of 52 late-type (Sa-Im-BCD) galaxies, combined with spectral evolutionary synthesis models show that the IMF of galaxies is consistent with a Salpeter IMF of slope $\\alpha$ = 2.35 and an upper mass cut off of about 80 ${\\rm M\\odot}$, in the mass range $m$ $>$ 2 ${\\rm M\\odot}$. Galaxies of different morphology and luminosity seem to have similar IMFs.  \\noindent 2) the birthrate parameter $b$, which gives the present-to-past SFR ratio ($SFR/<SFR>_{past}$), is strongly correlated with the total mass of galaxies as traced by their near-IR luminosity; the relationship between $b$ and the morphological type is weaker than with $L_H$.  \\noindent 3) the total gas reservoir per unit mass anti-correlates with the total mass of galaxies; as for $b$, the relationships between $M_{gas}$/$L_H$ and the morphological type are more dispersed than that with $L_H$. Low mass, dwarf galaxies, which are generally pure exponential discs, have a higher gas content per unit mass and present to past SFRs than early-type, massive spirals.  \\noindent 4) the birthrate parameter $b$ correlates with the total gas content of galaxies.  \\noindent 5) the SFE and the time scale for gas depletion do not strongly correlate with properties of galaxies such as the luminosity and/or the morphological type. If limited to pure discs, however, a trend between SFE and the total galaxy mass is observed. The time scale for gas depletion is $\\sim$ 10 Gyrs for normal, unperturbed late-type galaxies.  Point 1 is in agreement with the universality of the IMF in galaxies found in the high-mass range of normal galaxies (Scalo 1986, 1998; Massey 1998). Recent results based on star counts in OB associations in the Magellanic clouds seem to indicate that the IMF for massive stars is independent of metallicity (Hill et al. 1994). No systematic differences have been observed in the IMFs of massive stars associations in the Milky Way, LMC and SMC (Hill et al. 1994; Massey et al. 1995; Hunter et al. 1996c, 1997), M31 (Hunter et al. 1996a) and M33 (Hunter et al. 1996b); in all these nearby galaxies the IMF has a slope consistent with $\\alpha$=2.5 for masses $\\geq$ 1 ${\\rm M\\odot}$.   Points 2 and 3 have strong implications on the models of galaxy formation and evolution. The importance of the $b$ parameter resides on the fact that it gives directly an idea of the history of star formation of galaxies. In galaxies with a very small value of $b$ (generally early-type, massive galaxies) most of the stars have been formed at early epochs, and the present rate of star formation is lower than the past one. This is consistent with the idea that a rapid collapse might have induced a strong starburst, which efficiently transformed most of the gas into stars. The lack of gas at the present epoch makes these galaxies quiescent. Conversely, in objects with $b$ $\\sim$ 1 (late-type, dwarf, low mass galaxies) the gas is presently transformed into stars at the same rate than in the past. These observational evidences are consistent with the model of galaxy formation discussed by Sandage (1986), who proposed that galaxies are coeval systems, formed from the collapse of a primordial gas cloud, with a collapse time-scale depending on angular momentum. In elliptical galaxies the collapse was efficient enough to transform most of the gas into stars within a few 10$^8$ yr. In late-type systems the initial SFR was comparable with the present one (few solar masses per year), so that $\\sim$ 10 Gyrs after the formation of the primeval galaxy a large fraction of the gas is still available to feed new stars. This idea was recently rediscussed by Gavazzi et al. (1996a), Gavazzi \\& Scodeggio (1996), and Gavazzi et al. (1998). They observed that the $U-B$, $B-V$, $UV-B$ and $B-H$ colour indices and the H$\\alpha$ equivalent width depend strongly on the galaxy mass. They interpreted these observational evidences as an indication that the total mass contributes to the regulation of the process of collapse of primeval gas clouds. This idea is also supported by the recent numerical simulation by Noguchi (1999), indicating that the bulge-to-disc ratio seems primarily regulated by the total mass of the galaxy. A similar result was obtaind by Boissier \\& Prantzos (2000) and Boissier et al. (2000), whose chemo-spectrophotometric models of galaxy evolution reproduce several observed properties of disc galaxies only when the gas infall (and thus the star formation activity) is regulated by the total mass of the galaxy.  In the assumption that galaxies are coeval, and that they evolved as closed boxes, Gavazzi \\& Scodeggio (1996) and Gavazzi et al. (1998) were able to reproduce the relationships between the colour indices, the H$\\alpha$ equivalent widths and the $H$ luminosities, using the population synthesis models of Bruzual \\& Charlot (1993) and Kennicutt et al. (1994) and adopting a star formation law of the type:  \\begin{equation} {SFR(t) = SFR_o e^{(-t/\\tau)}    ~~~~~~~~ {\\rm M\\odot/yr}} \\end{equation}  \\noindent with a Salpeter IMF, where $t$ is the age of the galaxy, $\\tau$ is the time scale for collapse, and $SFR_o$ is the rate of star formation at the time of the galaxy formation ($t$=0). They found small values of $\\tau$ ($\\tau$ $\\sim$ 0.5 Gyr) for giant discs and large ($\\tau$ $\\sim$ 10 Gyr) for low mass objects.   The observational results of the present work are consistent with this simple picture. Let us assume that galaxies evolve as closed boxes following a star formation law of the type described in eq. (8), but assuming a more realistic age of $t_o$ = 12 Gyrs.   $M_{star}(t)$ at a given time $t$ is given by:  \\begin{equation} {M_{star} (t)= (1-R) \\int_{0}^{t} SFR_o e^{(-t'/\\tau)} dt' = (1-R) SFR_o \\tau [1 - e^{(-t/\\tau)}]} \\end{equation}  \\noindent \\footnote{as discussed in sect. 4.1 $R$ can be taken as a constant for $t$ $>$ 1 Gyr} that can be directly measured from the $H$ band luminosity as described in sect. 4.1:  \\begin{equation} {M_{star} (t)= L_H(t) DM_{cont}(t) [M_{tot}/L_H(t)]} \\end{equation}  \\noindent given the fact that the $H$ luminosity traces the total dynamical mass of galaxies. $SFR_o$ gives the number of stars formed per year at the formation of the galaxy. Its dependence on the total mass implies that the process of collapse is not scale-free. To simplify the formalism, we define $X(t) = DM_{cont}(t) (M_{tot}/L_H)$. As discussed in sec. 3.2, $X(t)$=0.5 $\\times$ 4.6 at $t=t_o$=12 Gyr.  \\noindent Combining eq. (5) with eq. (8) and (9) we can directly estimate the birthrate parameter $b$ as a function of $t$ and $\\tau$:   \\begin{equation} {b(t)= \\frac{t~e^{(-t/\\tau)}}{\\tau(1-e^{(-t/\\tau)})}} \\end{equation}  \\noindent The values of $b$ obtained from eq. (11) for various values of $\\tau$ and assuming $t$=$t_o$=12 Gyrs are consistent with those obtained from the population synthesis models of Kennicutt et al. (1994).    \\noindent If $M_{tot}$ is the total mass of a galaxy (constant with time in a closed box scenario), we can assume that:  \\begin{equation} {M_{tot} = M_{gas}(t) + M_{star}(t) + M_{DM}(t) ~~~~~~~~\\forall t} \\end{equation}  \\noindent where $M_{gas}$, $M_{star}$ and $M_{DM}$ are respectively the total gas, star and dark matter masses at a time $t$. If we assume that all the gas will be in the future transformed into stars, for $t$ $\\rightarrow$ $\\infty$ $M_{gas}(t)$ $\\rightarrow$ 0, while $M_{star}(t)$ = 0 for $t$=0, thus $M_{star}({\\infty})$ = $M_{gas}(0)$ = $M_{tot}$ - $M_{DM}(0)$ = $SFR_o$$\\tau$ (assuming that all the returned gas is used to form stars). If we assume that $M_{DM}$ does not change with time, $M_{gas}(t)$ can be determined from relation (12).  \\noindent By substituting $SFR_o$$\\tau$ to $M_{tot}$ - $M_{DM}$, eq. (12) becomes:  \\begin{equation} {M_{gas}(t) = SFR_o \\tau e^{(-t/\\tau)}} \\end{equation}  \\noindent and  \\begin{equation} {M_{gas}(t)/L_H(t) = \\frac {X(t)}{(1-R)}\\frac{e^{(-t/\\tau)}}{(1 - e^{(-t/\\tau)})}} \\end{equation}  \\noindent in the assumption that $R$ does not change with time.   At the same time this simple model predicts that the SFE defined in eq. (6) should be:  \\begin{equation} {SFE = \\tau^{-1}} \\end{equation}  \\noindent thus independent on $t$. This means that each galaxy should have had an efficiency in transforming gas into stars which depends only on the $\\tau$ parameter, but similar at any epoch.   This analytic representation of all the observed variables used in this work ($SFR$, $b$, $M_{gas}$, $L_H$, $SFE$) can be done once a star formation law such as that given in eq. (8) is assumed \\footnote{We have repeated this exercize replacing the exponential star formation law (eq. 8) with a ``delayed exponential'' law, such as the one proposed by Sandage (1986). This law better accounts for cases with increasing SFR with time, as observed in some irregular galaxies such as IZw18 (Kunth \\& Ostlin 2000). We found results qualitatively consistent with the exponential case, thus we avoid to expand the ``delayed exponential'' case in full details}.    The analytical model can be compared to the observables only once $SFR_o$ is known, $L_H$ being a function of $SFR_o$; the other variables ($b$, $M_{gas}$/$L_H$, $SFE$), being normalized entities, do not depend on $SFR_o$. This parameter, which gives the rate of star formation at the time of the galaxy formation ($t$=0), is clearly dependent on the mass of the galaxy ($SFR_o$ $\\times$ $\\tau$ = $M_{star}$ + $M_{gas}$).   Following Gavazzi \\& Scodeggio (1996), if we assume a simple empirical relationship between the $H$ luminosity and the exponential decay time scale for star formation of the type:  \\begin{equation} {log L_H = a \\times log \\tau + c} \\end{equation}  \\noindent we can derive from eq. (9) an (10) a semi-analytical relationship between $SFR_o$ and $\\tau$:  \\begin{equation} {SFR_o = \\frac {X(t_o) \\tau^a 10^c}{(1-R) \\tau e^{(-12/\\tau)}}} \\end{equation}  \\noindent Once $SFR_o$ is determined, we can compare the predictions of the analytical models with the relationships previously discussed between $b$, $M_{gas}$/$L_H$ or $SFE$ and the $H$ band luminosity (Fig 7 a,b,d). The model prediction (dotted line) can be directly compared to the observations for the $M_{gas}$/$L_H$ vs. $b$ relationship (Fig. 7 c), without any assumption on the relationship between $\\tau$ and $L_H$, both being normalized entities thus independent of $SFR_o$. At the same time the obvious observed scaling relationships between $SFR$, $M_{gas}$ and $L_H$ can be compared to the models prediction (Fig. 7 e,f) and used, together with the normalized entities, to constrain the free parameters $a$ and $c$. An excellent fit to the data is obtained for $a$ = -2.5 and $c$ = 12 once $\\tau$ is expressed in Gyrs consistent with  Gavazzi \\& Scodeggio (1996; $a$ = -2.5; $c$=11.12).  In conclusion, the observed relationships between optical, near-IR, UV colours, H$\\alpha$ E.W.s, the birthrate parameter $b$, the total gas mass and the $H$ luminosity can, at least to the first order, be reproduced in the simple assumption that galaxies are coeval ($\\sim$ 12 Gyrs), that they evolved as closed boxes following an exponentially declining star formation law, with a decay time scale depending on the mass of the primeval cloud, small ($\\tau$ $\\sim$ 0.5 Gyrs) for massive objects ($L_H$ = 10$^{12}$ $L_{H\\odot}$), and large ($\\tau$ $\\sim$ 10 Gyrs) for dwarf systems ($L_H$ = 10$^{9}$$L_{H\\odot}$).   If the models here discussed are realistic, we conclude that the strong relationship observed between various star formation tracers, such as the UV flux (Donas et al. 1990; Deharveng et al. 1994; Buat 1992; Boselli 1994), the H$\\alpha$ flux (Kennicutt 1989; 1998; Scodeggio \\& Gavazzi 1993 and references therein) and the HI or total gas surface density in unresolved galaxies is not due to a local relationship between the gas column density and the process of star formation, known as the Schmidt law, but it is simply a consequence of the closed-box evolution. From eq. (8) and (13) we expect in fact that the integrated SFR and the total gas content decrease with time with a similar decline. This conclusion is corroborated by the fact that in \"normal\", unperturbed galaxies the HI gas (which  constitutes $\\sim$ 90\\% of the total gas reservoir; Boselli et al. 1997b) is found in a flat disc of diameter about twice as large as the optical one (Cayatte et al. 1990). This gas is not expected to contribute directly to the star formation since i) up to 75\\% of the HI reservoir is located outside the optical disc where the star formation takes place; ii) the HI has to go through the molecular phase before forming stars. Thus an infall of gas from the outer HI disc to the optical disc must be present. Spiral galaxies of all morphological types and luminosity still have enough gas to continue their present day star formation activity for $\\sim$ 10 Gyrs. This time is significantely longer than previous estimates because we assume an higher contribution of the recycled gas (see Kennicutt et al. 1994) and because the HI-deficient cluster galaxies are excluded from the analysis (Boselli 1998). ",
        "conclusions": ""
    },
    "0011/astro-ph0011220_arXiv.txt": {
        "abstract": "We present a proper motion and CCD photometric study of stars in the distant halo globular cluster Palomar 13. The absolute proper motion of Pal 13 with respect to the background galaxies, derived from moderate scale photographic plates separated by a 40-year baseline, is $(\\mu_{\\alpha cos \\delta}, \\mu_{\\delta}) = (+2.30, +0.27) \\pm (0.26, 0.25)$ milliarc-seconds per year.  The resultant total space velocity (315 km s$^{-1}$) implies that Pal 13 is in the inner part of its orbit near perigalacticon. Orbital integration reveals the cluster to possess an inclined, very eccentric, retrograde orbit.  These data confirm that Pal 13 is a paradigm ``young halo\" globular cluster.  The derived proper motions for cluster stars are used to produce membership probabilities and a cleaned CCD UBV catalogue for Pal 13.  With this data set we have made small revisions to Pal 13's distance, metallicity, position and light profile.  The membership of four previously reported RR Lyrae variables and a proportionally large group of blue straggler stars are confirmed.  As expected, the blue stragglers are centrally concentrated.  The small size of this cluster, combined with the shape of its light profile, which shows a clear departure from a classical King function beyond the tidal radius, suggests that Pal 13 is in the final throes of destruction. This could explain the large blue straggler specific frequency, as destructive processes would preferentially strip less massive stars. ",
        "introduction": "The globular cluster Palomar 13 is one of a number of faint, low mass, sparse globular clusters identified from the Palomar Observatory Sky Survey (Abel 1955; Wilson 1955).  The first study of this cluster by Ciatti et al. (1965) identified four RR Lyrae stars in the cluster and derived a distance modulus of $(m-M)_{V} =17.11$.  At $(\\alpha, \\delta)_{1950}=(23:04.2, +12:28)$, or $(l,b)=(87.1^{\\circ}, -42.7^{\\circ}$), this distance modulus places Pal 13 at a Galactocentric radius of 27.2 kpc, well into the halo of the Milky Way.  Further information on Pal 13 did not appear in the literature for over a decade when a flurry of spectroscopic studies of Pal 13 revealed it be moderately metal poor.  Spectroscopic [Fe/H] were measured at $-1.9 \\pm 0.4$ (Canterna \\& Schommer 1978), $-1.75 \\pm 0.6$ (McClure \\& Hesser 1981), $-1.67 \\pm 0.15$ (Zinn \\& Diaz 1982) and $-1.9 \\pm 0.1$ (Friel et al. 1982).  In the most comprehensive study to date, Ortolani et al. (1985, hereafter ORS) reexamined the 200-inch Palomar plates that were the basis for the Ciatti et al. study (as well as the present astrometric work).  ORS confirmed the distance modulus, determined a metallicity between [Fe/H]=-1 and -1.5, and found that Pal 13 was the faintest and smallest globular cluster measured at that time\\footnote{Harris (1996) currently lists Pal 13 as the fifth smallest globular cluster based upon its luminosity.}.  They claimed that Pal 13's luminosity function is similar to that of the bright globular cluster M3, but scaled down by a factor of $\\sim 60$.  From this, ORS estimated an absolute magnitude of $M_{V}$=-3.36 and a mass of approximately 3 x $10^3$ solar masses with an assumed (M/L) ratio of 1.6.  Recently, the CCD photometry of Borissova et al. (1997, hereafter B97) revealed Pal 13 to have an age of approximately 12 Gyr based on the fitting of isochrones from Proffitt \\& VandenBerg (1991) and Bergbusch \\& VandenBerg (1992) - suggesting that it is representative of the ``young halo\" population of globular clusters.  In addition, B97 reported a number of possible blue straggler stars, a Zinn-West (1984) metallicity of [Fe/H]=-1.52 and confirmed the Ciatti et al. distance modulus.  We have undertaken a comprehensive photometric and astrometric study of Pal 13 to improve understanding of this object.  The present work updates and supercedes our previous report (Cudworth et al. 1993).  Our efforts have been directed toward three goals:  (1) The measurement of Pal 13's absolute proper motion and subsequent determination of its true space velocity.  This will improve our knowledge of Pal 13's orbit, the Galactic mass enclosed by that orbit and provide insights into its dynamical history.  (2) The production of an improved UBV CCD photometric catalogue of the Pal 13 field.  (3) The use of proper motions to separate true members of Pal 13 from the field star population.  A cleaned color-magnitude diagram (CMD) will help to verify the membership of blue stragglers and RR Lyrae variables.  It will also improve our ability to compare isochrones to confirm Pal 13's age and metallicity. This is particularly important for Pal 13, which is a sparse cluster and does not show a strong contrast against the field stars except in its core.  This is the fifth object and third globular cluster in our program (Majewski \\& Cudworth 1993) of measuring absolute proper motions of distant globular clusters and dwarf spheroidal galaxies using faint galaxies and QSO's to fix the proper motion zero point (Schweitzer et al. 1995; Dinescu et al. 2000; Schweitzer et al. 2000; Cudworth et al. 2000).  \\S2 describes our observations and \\S3 our reduction methods.  Our photometric and astrometric results are presented in \\S4 and \\S5 respectively and the implications of our work are discussed in \\S6. ",
        "conclusions": "\\subsection{The Mass of the Milky Way}  Given the high space velocity of Pal 13, it is natural to ask whether the cluster is bound to the Galaxy.  The minimum Galactic mass needed to bind Pal 13 can be derived from its Galactocentric radius and velocity by a simple escape velocity calculation.  Using the formulation $M=v_{esc}^2r/2G$, an approximate mass of 3 x $10^{11} M_{\\sun}$ interior to its present position would keep Pal 13 bound.  Estimates of the Milky Way mass have a large degree of uncertainty.  Zaritsky et al. (1989) estimated the mass of the Galaxy to be 9-12 x $10^{11} M_{\\sun}$ if Leo I is bound, 4 x $10^{11} M_{\\sun}$ if not.  Cudworth (1990) used local stars to place a minimum Galactic mass at 4x$10^{11} M_{\\sun}$.  Using the velocity distribution of outer halo objects as revealed by proper motions and radial velocities, Kulessa \\& Lynden-Bell (1992) have estimated a Galactic mass as high as $10^{12} M_{\\sun}$ while Kochanek (1996) estimated a mass of 3.9-5.1 x $10^{11} M_{\\sun}$ inside 50 kpc, depending on the status of Leo I.  Most recently, Wilkinson \\& Evans (1999) used this method to estimate a Galactic mass of 5 x $10^{11} M_{\\sun}$ inside 50 kpc independent of Leo I.  The proper motions of Pal 13 and other distant halo objects could precipitate substantial improvements upon these mass estimates (c.f. Wilkinson \\& Evans 1999) and this is one of the ultimate goals of our astrometric survey of outer halo objects.  The current estimates of the Milky Way mass are all high enough to bind Pal 13.  \\subsection{Orbit and Dynamical History of Pal 13}  The well-established position and distance of Pal 13 in combination with our new space velocity can be used to predict the cluster's future orbit or outline its past orbit.  For this purpose, we have used the orbital integration tool and model of Johnston et al. (1995, hereafter JSH95).  Figure 8 shows the orbit of Pal 13 integrated 10 Gyr into the future.  The rosette pattern that results from its present velocity is fairly typical of outer halo globular clusters (c.f. Dinescu et al. 1999, hereafter D99).  Table 7 lists the calculated orbital parameters of Pal 13, based on the true space velocity presented in this paper and the JSH95 potential.  Orbital elements were calculated as in D99 with the small changes that inclination angle ($\\Psi$) is calculated at apogalacticon and the period is the perigalactic period as opposed to the azimuthal period.  As expected from the high velocity, Pal 13 spends most of its orbital time in the distant halo.  It is presently only 70 Myr past perigalacticon and spends only 12\\% of its 1.1 Gyr orbit this close to the Galactic Center.  The non-color-corrected proper motion of Pal 13 produces a much more radical orbit with apogalacticon at 185 kpc and a 2.6 Gyr orbital period.  The high velocity given by the uncorrected proper motion is so extreme that any orbital integration is uniquely sensitive to model parameters in a way that other cluster motion integrations are not.  Therefore, the associated orbital parameters are very uncertain because the JSH95 model uses an infinite logarithmic potential for the Galactic halo.  Under these circumstances, the apogalacticon and eccentricity would represent only lower bounds on Pal 13's orbit.  In either the color-corrected or uncorrected case, Pal 13 is on a highly eccentric, retrograde, inclined orbit and spends the vast majority of its time in the distant reaches of the Galactic halo.  This point is discussed further in \\S6.5.  Lynden-Bell \\& Lynden-Bell (1995) proposed two dynamical associations of Pal 13 with dwarf spheroidal galaxies and other globular clusters and made predictions for its proper motion supposing such association.  Neither of these proper motions is close to the value we have derived for Pal 13.  Majewski (1994) suggested a possible association with the Magellanic stream.  Figure 9 shows the orbital poles of Pal 13 and various proposed ``dynamical families\", using the methods detailed in Palma et al. (2000).  We find that Pal 13's orbital pole is significantly removed from that of any previously proposed family, including the Magellanic Stream, the Fornax-Leo-Sculptor Stream and the Sagittarius dSph.  Pal 13's association with some of the presently known Galactic satellites can also be ruled out by comparing their dynamical properties and histories.  Pal 13 has a very high eccentricity.  If Pal 13 had been tidally stripped from a parent object, that interaction would leave Pal 13 on an orbit that is 10-30\\% more or less energetic than its parent object (Johnston 1998).  As this stripping would happen at or near the pericenter of the parent object's orbit, this would result in a less eccentric orbit for Pal 13 if its on a less energetic orbit, more eccentric if on a more energetic orbit.  We have compared Pal 13's dynamical properties to the five satellite galaxies for which proper motions have been measured.  These are Draco (Scholz \\& Irwin 1994), Ursa Minor (Schweitzer et al. 2000), Sculptor (Schweitzer et al. 1995), Sagittarius (using the HST proper motion of Ibata et al. 2000, which is similar to the ground-based measure of Irwin et al. 1996) and the LMC (an average of the motions from Jones et al. 1994, Kroupa et al. 1994 and Kroupa \\& Bastian 1997) \\footnote{We note that a recent LMC proper motion has been published by Anguita et al. (2000), where the proper motion has been referenced to three quasars.  We have not included this result primarily because it is substantially different from the two previous independent measures and implies a space velocity in excess of the escape velocity of the Milky Way.  Our own experience (Majewski 1992) using quasars as part of an extra-galactic reference frame suggests problems in the Anguita et al. measures resulting from differences in the spectral energy distributions of stars and quasars.}. We have combined these proper motions with radial velocities, distances and positions from Mateo (1998) and references therein to derive orbits for the satellite galaxies, and then compared these orbits to that of Pal 13.  Table 8 lists the proper motions, orbital energy, average total angular momentum, perigalacticon, apogalacticon and eccentricity for Pal 13 and the dSph galaxies with measured proper motions.  We find that we can rule out all of the dwarf galaxies with measured proper motions from an association with Pal 13 under the scenario of a simple Milky Way-satellite evolution. Draco's proper motion implies a velocity too large for it to be bound to the Milky Way.  The JSH95 model has an infinite halo that will eventually pull Draco back into the Galaxy.  However, such a halo is clearly not appropriate if Draco's velocity is as extreme as suggested by its proper motion.  The large uncertainties do include a number of bound orbits but do not allow a good dynamical comparison to Pal 13.  The LMC, Sculptor and Ursa Minor have more energetic orbits than Pal 13 but smaller eccentricities.  Sagittarius has 24\\% less energy than Pal 13 and a smaller eccentricity.   However, Sagittarius has an angular momentum approximately 40\\% smaller than that of Pal 13, a difference too substantial for a common origin.  These energy and orbital comparisons are, of course, only valid in the case of secular evolution of a two-body Milky Way-dSph interaction.  They do not rule out three-body-encounters, e.g. the collision of two satellites in the Galactic halo as explored by Zhao (1998).  \\subsection{Comparing Two Measures of Perigalacticon}  Because Pal 13 has an eccentric orbit, the cluster would seem to be an interesting object with which to test the assertion that the King profile limiting radius ($r_K$) reflects the tidal radius of the cluster at perigalacticon  (King 1962; Innanen et al. 1983).  Such an argument was used by ORS to speculate that Pal 13 was near perigalacticon.  Oh \\& Lin (1992) tested orbits of $e=0.5$ using a Fokker-Planck code and found no orbital phase dependence for the limiting radius, which indicates that a cluster's limiting radius should always reflect its perigalactic tidal radius.  Pal 13 would present a unique observational test of this scenario as the cluster will be in the distant halo for a period longer than its relaxation time (3 x 10$^8$ years, Gnedin \\& Ostriker 1997) and could expand to a larger radius as its tidal radius expands with declining Galactic potential.  Using the formulation of Innanen et al., our revision of $r_K$, a Pal 13 mass of $5 \\times 10^3 M_{\\sun}$ (based on a more modern M/L ratio to 3.0, Gnedin \\& Ostriker 1997) and a Galactic mass of $5 \\times 10^{11} M_{\\sun}$ we estimate the perigalacticon of Pal 13 to be 26 $\\pm$ 1 kpc, which is significantly larger than the perigalacticon calculated above from the cluster space motion (11.2 kpc), and is within error bars of the cluster's {\\it present} distance from the Galaxy.  The inconsistency between the kinematic and tidal estimates of Pal 13's perigalacticon is notable and perhaps cautions against assuming that the limiting radius and perigalactic tidal radius are identical.  \\subsection{The Fate of Pal 13}  Pal 13's extremely small size is either the result of having been formed as a small globular cluster or having been reduced in size by internal or external forces.  Gnedin \\& Ostriker (1997) estimate that the only destructive process that will have major effect upon this particular cluster is evaporation, a mechanism that will completely destroy the cluster within the next Hubble time\\footnote{The possibility of disruption by a dark halo of black holes has been addressed by Moore (1993) with a specific analysis of Pal 13.  The conclusion was that Pal 13 does not show evidence of having been disrupted by such encounters.}. However, their evaluation assumed a proper motion for Pal 13 consistent with a kinematical model.  The fact that Pal 13 passes much closer to the Milky Way than previously suspected and, on closest approach, will not be able to retain the stars in the outer edge of its King profile, makes it more vulnerable to disk and/or tidal shocking than anticipated in the Gnedin \\& Ostriker study.  The possibility of a more rapid destruction of Pal 13 seems likely.  Clues to substantial mass loss heretofore lie in Pal 13's present stellar population.  Equipartition of energy among cluster stars means that evaporation and tidal stripping preferentially occur for lower mass stars, while more massive stars, particularly binaries, will be left behind in the cluster.  Pal 13's CMD yields two indicators of a disproportionate binary content compared to ``normal clusters\".  The first is the apparent double subgiant branch, possibly made more evident in Pal 13 due to the preferential reduction of non-binary subgiant stars with time.  However, as we have noted, the stars in the ``double subgiant branch\" are not found to be centrally concentrated, as might be expected.  The second clue, of course, is the relatively substantial BSS population for a cluster of this mass and miniscule RGB content.  We have already shown that the BSS population is concentrated to the core showing that mass segregation has likely occurred.  If BSS are understood to be merged binaries then their high specific frequency in low mass, low concentration clusters like Pal 13 (c.f. Fusi-Pecci 1992) can be understood as the result of substantial tidal stripping.  Indeed, Pal 13 is almost a twin of the globular cluster E3 in BSS specific frequency, mass and concentration (van den Bergh et al. 1980).  van den Bergh et al. argue these traits to be strong evidence of the death of E3, and we believe these arguments apply to Pal 13 as well.  Given the apparent strong past evolution of Pal 13, it is likely that it is on the fast track to complete destruction by tidal forces.  Rapid evaporation would produce the mass segregation and preferential lingering of high mass stars, but would not explain the low concentration parameter.  Cluster evaporation elevates the concentration parameter preceding core collapse (c.f. Elson 1999).  Of course, we have evidence of substantial mass loss in Pal 13 now, as more than 10\\% of its members stars are outside the nominal limiting radius (Figure 7) which would be equal to the tidal radius at Pal 13's present distance.  These stars are too far away from Pal 13's core to be bound to the cluster at its present Galactocentric distance.  In addition, they appear to follow an $r^{-\\alpha}$ distribution as predited by Johnston et al. (1999, 2000) and confirmed in the case of Carina by Majewski (2000b).  However, the stars are too few and too far apart for meaningful profile fitting.  \\subsection{Implications for Galactic Formation}  The globular cluster system of the Milky Way has long been thought to hold clues to the formation of the Galaxy (c.f. Majewski 1999).  It was the lack of an abundance gradient in the halo globular clusters that led to the proposal that part of the Galaxy was formed from diverse fragments (Searle \\& Zinn 1978).  This has been more recently refined (Zinn 1985; Zinn 1993; D99) into a scenario in which the Milky Way clusters are divided into three populations - a disk population of metal-rich clusters, and two halo populations of metal poor clusters.  The ``young halo\" clusters would represent an accreted portion of the Galaxy while the ``old halo\" population would represent the portion formed during the collapse that eventually formed the disk\\footnote{Zinn (1996) divides the former ``old halo\" population into ``metal-poor\" and ``blue horizontal branch\" populations.  However, as both were likely formed during the collapse of the Galaxy, the distinction between them is unimportant for the present discussion.  Pal 13's classification would be identical in both schemes.}.  The ``young halo\" clusters are characterized by younger age, a second parameter effect in their HB, a mean retrograde rotation, more eccentric orbits and higher $Z_{max}$ values. Pal 13 conforms almost perfectly to this paradigm.  While it has only five horizontal branch stars, Pal 13's metallicity implies that they would be very blue in the absence of a second parameter effect.  This is clearly not the case. Moreover, Pal 13 is a relatively young cluster (according to the MSTO fitting of B97), has a very high $Z_{max}$ and is on a retrograde orbit.  An object with this orbit would not be consistent with Zinn's old halo population unless it were extremely metal poor.  In the context of accretion, Pal 13 is unlikely to have been accreted from the nearby dwarf galaxies with measured proper motions (\\S6.2).  However, it could have been accreted from an undiscovered dwarf or one with an unmeasured proper motion.  Clearly, insight into the cluster population of the Milky Way will require the measurement of proper motions for more outer halo globular clusters and dwarf galaxies, which we are attempting in our ongoing survey. However, Pal 13 does pose an interesting question as it is in a very fortunate position.  It is intrinsically faint to the point of being almost unnoticeable on even deep exposures.  It is also near its closest point to the Milky Way - the point on its orbit where it will spend the least time.  Our orbital integration shows that Pal 13 spends only a small fraction of its time at its present distance or closer to the Galactic Center. We note that the faint halo globular cluster Pal 12 is also near its perigalactic point (Dinescu et al. 2000). It is impossible to determine if such detections are mere coincidence or reflective of a large population of faint distant halo globular clusters that are only detectable near perigalacticon.  How many more Pal 13's are in the outer halo, too faint and too far away to have been discovered?  The heavens have been thoroughly searched for such objects and the outer halo census includes such faint objects as Pal 4 ($M_V=-5.82, R_{\\sun}=100$ kpc),  Pal 3 ($M_V=-5.59, R_{\\sun}=90$ kpc), Eridanus ($M_V=-4.90, R_{\\sun}=80$ kpc), Pal 14 ($M_V=-4.68, R_{\\sun}=72$ kpc) and AM-1 ($M_V=-4.66, R_{\\sun}=129$ kpc) (Harris 1996 and references therein).  Any undiscovered cluster in the outer halo would have to be even fainter than these objects or at low Galactic latitude.  We note that nine clusters within 25 kpc of the Sun are fainter.  Would it be possible for a Pal 13-sized object to be missed at large radii?  Pal 13 is more than a magnitude fainter than the clusters listed above.  At a distance of 100 kpc, Pal 13 would have only have one or two members brighter than $V=20$ and only 15-16 members brighter than $V=21$.  Such an object would not have been detected by the POSS.  Thus, the possibility of a number of Pal 13-sized outer halo globular clusters remains open.  The Palomar globular clusters were discovered by the POSS, a deliberate and planned effort to search the heavens for such objects.  The ongoing SLOAN digital sky survey will extend our reach even further and hopefully give insight into whether Pal 13 and Pal 12 are simply rare objects at the tail of the Milky Way globular cluster mass distribution or the brightest, nearest representatives of a group of very sparse, distant halo clusters."
    },
    "0011/astro-ph0011199_arXiv.txt": {
        "abstract": "I summarize recent results from radio observations of high redshift star forming galaxies, discuss radio continuum emission as a measure of star formation rate, and consider future capabilities at cm to IR wavelengths. ",
        "introduction": "Source counts based on low frequency surveys of the sky with Jy sensitivity showed a significant departure from a Euclidean, non-evolving source population, indicating, for the first time, cosmic evolution in a source population. The source population entailed luminous radio galaxies, with spectral luminosities at 178 MHz: $\\rm P_{178} > 10^{32} ~erg~ s^{-1}~ Hz^{-1}$ \\cite{ryle,jauncey}. In these sources the synchrotron  radio emission is from high energy electrons accelerated in a relativistic jet emanating from the active galactic  nucleus (AGN)\\cite{bartel}.  Subsequent observations with sub-mJy sensitivity, starting with the WSRT and continuing with the VLA \\cite{windhorst}, revealed flattening of the  source counts below 5 mJy (Fig. 1). Windhorst et al. hypothesized that this flattening was due to a new population of sources, namely star forming galaxies with $\\rm P_{178}  < 10^{31} ~erg~ s^{-1}~ Hz^{-1}$. The radio emission in these sources is from relativistic electrons accelerated in  supernovae remnant shocks\\cite{condon}.  There has been a recent revival of deep radio surveys, motivated in large part as follow-up to deep optical, infrared, and (sub)mm surveys \\cite{richards,hopkins,garrett,owen}.  The frequency of choice is 1.4 GHz for these deep surveys, allowing for $\\mu$Jy sensitivity with arcsecond resolution and a wide field of view (FWHM = 30$'$). In the coming years the Expanded VLA (EVLA) will push to the sub-$\\mu$Jy level, while in the coming decades the Square Kilometer Array (SKA) will potentially probe nJy sources.  As pointed out in numerous papers in these proceedings (Bertoldi, Adelberger, Hughes, Sanders), radio observations play an important, complimentary role to observations at other wavelengths, in that: \\begin{itemize} \\item They are not plagued by extinction corrections.  \\item They provide a rough  estimate of source redshifts, in combination with (sub)mm observations \\cite{yun99,yun00}.  \\item Low order CO transitions redshift into the cm bands, revealing large reservoirs of less dense, cooler gas \\cite{papadopoulos,menten}. This topic will not be discussed herein.  \\item They provide arcsecond astrometry and imaging, thereby avoiding the confusion problems inherent in deep searches for optical counterparts of (sub)mm sources discovered in low resolution single dish bolometer array surveys. \\end{itemize}  \\noindent For example, at the optical limits of the HDF ($I_{AB} < 29$) one expects, by chance, three faint galaxies within the 6$''$ error circle of a typical SCUBA or MAMBO source\\cite{lilly}. At 1.4 GHz the source counts between 40 $\\mu$Jy and 1 mJy obey: $\\rm N(> S_{1.4}) = 2.2\\times10^{-6}~ S_{1.4}^{-1.4}$~ arcsec$^{-2}$, with S$_{1.4}$ in mJy\\cite{richards}. The number of  spurious $\\rm S_{1.4} \\ge 22 \\mu$Jy sources within the error circle is only 0.02.  \\begin{figure} \\vskip -0.1 in \\hspace*{-0.2in} \\psfig{figure=carilli_fig1.ps,width=13cm} \\vskip -0.3in \\caption{Radio source counts at 1.4 GHz (adopted from \\cite{hopkins2}). } \\end{figure}  Since deep radio images are usually  follow-up observations of deep fields at other wavebands, there has been very rapid progress in determining the nature of the parent galaxies of the $\\mu$Jy radio source population \\cite{richards,haarsma}. This short summary will focus on recent results on the $\\mu$Jy radio source population, emphasizing the unique information about high redshift star forming galaxies coming from radio observations. I will also discuss the radio-to-far infrared (FIR) correlation, re-deriving star formation rates based on this correlation, and conclude with a short discussion of future instrumentation.  Richards \\cite{richards} finds that 75$\\%$ of the S$_{1.4} \\ge 40\\mu$Jy radio sources are identified to $I < 24$, with a median of $I = 22.1$. Interestingly, he also finds that 25$\\%$ of the sources are unidentified to $I > 25$. And perhaps more interestingly, a number of groups \\cite{richards,barger99,bertoldi1,chapman} find that (sub)mm observations of these $\\mu$Jy radio sources with  faint (or absent) optical counterparts results in a $> 40\\%$ detection rate at mJy levels. This suggests that the optically faint, $\\mu$Jy radio sources are equivalent to the mJy (sub)mm source population.  Haarsma et al. \\cite{haarsma} have used an extensive spectroscopic and photometric redshift analysis to determine the redshift distribution of the $\\mu$Jy radio source population. They find that (roughly): \\begin{itemize} \\item 50$\\%$ of the sources are spirals, or irregular galaxies, at $z < 1$.  \\item 25$\\%$ are ellipticals, presumably low luminosity AGN, at $0.3 < z < 1.5$.  \\item 25$\\%$ are optically faint (or absent), and red. \\end{itemize}  \\noindent They propose that these later sources are likely to be high redshift ($1.5 < z < 4$), dust obscured  starbursts. This idea is consistent with the SCUBA and MAMBO results discussed above \\cite{richards,barger99,bertoldi1,chapman}.  The angular size distribution of the $\\mu$Jy radio source population remains a point of debate. WSRT observations of the HDF at 1.4 GHz, 15$''$ resolution, to 8$\\mu$Jy rms detect a number of sources not detected in the VLA survey at similar sensitivity but with 1$''$ resolution \\cite{garrett}.  This would suggest a significant population of sources larger than 1$''$. However, confusion is a serious issue at this sensitivity level at 15$''$ resolution. Combined MERLIN+VLA observations of the HDF suggest that most of the sources have angular sizes between 0.7$''$ and 2$''$, with a median of 1.4$''$ (Muxlow et al. in prep). On the other hand, VLA observations of the cluster A2125 suggest that most of the sources are unresolved, with upper limits of typically 1$''$\\cite{owen}.  If the sources are indeed 1$''$ in size, this presents a significant problem for the SKA, since the sky will become naturally confusion limited at the few nJy level, independent of the resolution of the instrument \\cite{hopkins}, well above the sensitivity of the SKA.  Richards \\cite{richards} finds that the mean spectral index for a 1.4 GHz selected sample of sources is --0.8, typical of star forming galaxies. Not surprisingly, an 8 GHz selected sample shows a flatter mean spectral index of --0.4.  Barger et al. \\cite{barger99} find that only 15$\\%$ of the radio sources with S$_{1.4} \\ge 25 \\mu$Jy in the Hawaii Deep Field are X-ray sources with 2 to 10 keV fluxes of: $\\rm I_x \\ge 1\\times10^{-15}$ erg s$^{-1}$ cm$^{-2}$. This result is  consistent with the idea that the majority of $\\mu$Jy radio sources are star forming galaxies.   One problem with deep radio surveys is the limited area covered, such that cosmic variance can lead to substantial differences between counts derived for different fields. For instance, the counts at the 100 $\\mu$Jy level in the Phoenix Deep Field \\cite{hopkins} are a factor two higher than those in the Hubble Deep Field \\cite{richards}. This variance can be seen as the increased scatter below 1 mJy in Fig. 1. ",
        "conclusions": ""
    },
    "0011/astro-ph0011366_arXiv.txt": {
        "abstract": "We used the High Energy Transmission Grating Spectrometer (HETGS) onboard the Chandra X-ray Observatory to perform two observations, separated by three weeks, of the Orion Trapezium region. The zeroth order images on the Advanced CCD Imaging Spectrometer (ACIS) provide spatial resolution of 0.5\" and moderate energy resolution. Within a 160\"x140\" region around the Orion Trapezium we resolve 111 X-ray sources with luminosities between 7$\\times 10^{28}$ \\ergsec and 2$\\times 10^{32}$ \\ergsec. We do not detect any diffuse emission. All but six sources are identified. From spectral fits of the three brightest stars in the Trapezium we determine the line of sight column density to be N$_{H}$ = 1.93$\\pm0.29 \\times 10^{21}$ cm$^{-2}$. Many sources appear much more heavily absorbed, with N$_{H}$ in the range of $10^{22}$ to $10^{23}$ cm$^{-2}$. A large fraction of sources also show excursions in luminosity by more than a factor 5 on timescales $>$50 ks; many are only detected in one of the observations.  The main objective of this paper is to study the Orion Trapezium and its close vicinity. All five Trapezium stars are bright in X-rays, with $\\theta^1$ Ori C accounting for about 60$\\%$ of the total luminosity of the Trapezium. The CCD spectra of the three very early type members can be fit with a two-temperature thermal spectrum with a soft component of kT $\\sim$ 0.8 keV and a hard component of kT $\\sim$ 2 to 3 keV. $\\theta^1$ Ori B is an order of magnitude fainter than $\\theta^1$ Ori E and shows only a hard spectrum of kT $\\sim$ 3 keV. $\\theta^1$ Ori D is another order of magnitude fainter than $\\theta^1$ Ori B, with only a  kT $\\sim$ 0.7 keV component. We discuss these results in the context of stellar wind models.  We detect eight additional, mostly variable X-ray sources in the close vicinity of the Trapezium. They are identified with thermal and non-thermal radio sources, as well as infrared and optical stars. Five of these X-rays sources are identified with proplyds and we argue that the X-ray emission originates from class I, II and III protostars at the cores of the proplyds. ",
        "introduction": "The Trapezium Cluster of the Orion Nebula is one of the best studied star forming regions. At a distance of 440 pc, the Trapezium Cluster is embedded at the near rim of the giant dark molecular cloud L1641, lying just behind the Orion Nebula (NGC 1976). It is the core of the Orion Nebula Cluster (ONC), a loose association of more than  5000 mostly very young ($<$ 1-2 Myr) pre-main sequence stars of a wide range of stellar mass ($\\sim 0.1 - 50$ \\Msun) within a radius of about 1.3 pc (see Hillenbrand 1997 and references therein). A density of $\\sim 5 \\times 10^3$ stars per pc$^3$ within the central 0.1 pc makes the Trapezium Cluster the most dense stellar clustering known in our Galaxy (McCaughrean $\\&$ Stauffer 1994) with about 3000 \\Msun pc$^{-3}$.  The Orion Trapezium was discovered by Trumpler (1931) and its main stellar component, $\\theta^1$ Ori C, was identified as the main source of excitation of the Orion Nebula. Many observations of the Trapezium have been made at infrared wavelengths, mainly the K- and H-bands at 2.20 $\\mu$m and 1.65 $\\mu$m respectively. $\\theta^1$ Ori A  - E are the brightest stars in the cluster, optically and at infrared wavelengths, and  are early type  O or B-stars. Three members of the Trapezium are multiple systems, some with separations much smaller than 90 AU (Petr et al. 1986). About 80 $\\%$ of the stars are younger than 1 Myr with a median age of 3$\\times 10^5 $ yr. The present star formation rate is estimated at $\\sim 10^{-4}$ \\Msun yr$^{-1}$ (Hillenbrand 1997).  The core of the Orion Nebula is one of the prime areas to observe and study very young stellar objects (YSOs). Laques $\\&$ Vidal (1979) detected externally illuminated YSOs in the form of compact H$\\alpha$ brightness knots, which are illuminated by the two optically brightest sources, $\\theta^1$ Ori C and $\\theta^1$ Ori A. Garay et al. 1987 measured free-free radio continua from most of these sources. Many so-called 'proplyd' (externally illuminated protoplanetary disk) candidates have been discovered in the vicinity of the Trapezium (O'Dell et al. 1993), many of which correspond to radio and/or IR sources (Felli et al. 1993; McCaughrean and Stauffer 1994). Bally et al. (1998) studied over 40 proplyds with HST and discussed models for these illuminated YSOs involving a young star at the core with a circumstellar disk embedded in an outer wind shock region, as well as a wind driven tail pointing away from the illuminating star.  Because of its proximity, the Orion Trapezium came under early scrunity in X-rays and it was UHURU, the first X-ray satellite (Giacconi et al. 1971), that discovered X-rays from that region in 1972. The \\ein Observatory for the first time observed X-rays identified with main sequence O6 to B5 stars (Cassinelli et al. 1981). These early type stars do not have a stellar corona like late type stars, but exhibit a very hot stellar wind that ionizes its environment. Here X-rays are generally thought to originate from bow shocks surrounding radiatively driven blobs plowing through rarified wind material (Lucy and White 1980). Surprisingly, X-rays from some late-B and early A-type stars were discovered, where neither a corona nor an important stellar wind is thought to exist. The generally favored explanantion for these X-rays is that they originate not from the star itself but from an unresolved, otherwise invisible low-mass companion (Bergh\\\"ofer and Schmitt 1994). On the other hand the ONC hosts a large sample of late type pre-main sequence (PMS) stars. \\ein and \\ros established a large catalog of X-ray sources within an area of 5'x5', of which over 200 objects were identified as PMS stars with spectral types of F to late G (Ku $\\&$ Chanan 1979, Gagne et al. 1995, Geier et al. 1995). These stars generate X-rays through magnetic activity and a hot corona. It has also been suggested that the Trapezium Cluster itself features a faint diffuse X-ray component from ionized material in between the Trapezium stars (Ku $\\&$ Chanan 1979).  With the launch of ROSAT in 1990, individual X-ray sources were resolved in the Trapezium (Caillault et al.  1994; Gagne et al. 1995), however the question of a diffuse X-ray component within the Trapezium remained (Geier et al. 1995). The \\asca observatory launched in 1993 observed the Orion Trapezium cluster with its until then unprecedented spectral resolving capacity up to 10 keV. The results showed that a high temperature plasma with kT $\\sim 2-5$ keV exists in the Trapezium, and 2-temperature thin thermal plasmas are more likely than single temperatures (Yamauchi et al. 1996). This is in agreement with the observations of O-stars with \\asca (Corcoran et al. 1994), which obtained spectra of $\\delta$ Ori and $\\lambda$ Ori characterized by two-temperature models with additional absorption attributed to a warm wind. The hot components had temperatures of kT $\\sim$ 0.6 and 2 keV. The \\asca observations also did not observe the large flares one expects from late type T Tauri stars, but the spectra showed the existence of a variable hard spectral component associated with the Trapezium.  The physical origin of highly energetic emission from young stars is still uncertain. So far only empirical relationships exist between this emission and stellar age, size and mass of the star, as well as stellar rotation (Feigelson et al. 1993, Kastner et al. 1997). We currently classify evolutionary traces simply by observability in the radio to X-ray band. Class 0 sources are considered to be still deeply embedded and merely gravitationally collapsing cores and the spectral energy distribution peaks in the far-IR/submm range. Class I sources are accreting from a protostellar cloud and are optically obscured. The slope of these spectra is known to rise from near-IR through far-IR. Class II sources or classical T Tauri stars (CTTS) are strongly visible in the optical in the H$\\alpha$-line but also have a substantial IR excess. Class III sources or weak line T Tauri stars (WTTS) are optically bright with little or no IR excess; see the review by Feigelson $\\&$ Montmerle (1999) and references therein. \\asca recently detected giant X-ray flares with variable hard X-ray emission from YSOs (Kamata et al. 1997, Tsuboi et al. 1998, Yamauchi $\\&$ Kamimura 1999, Tsuboi et al. 2000). Spectral temperatures ranged from kT $\\sim$ 4 - 6 keV, in extreme cases up to $>$ 10 keV. A recent modeling of these X-ray flares from two class I protostars relates this emission to strong magnetic shearing and reconnection between the central star and the sourrounding protostellar accretion disk (Montmerle et al. 2000).  The next steps toward a better understanding of the astrophysical processes underlying high energy emission at early stages of stellar evolution can be made with the Chandra X-ray Observatory. The Orion Nebula Cluster was observed early in the Chandra mission by Garmire et al. 2000, who detected about 1000 resolved X-ray sources, the richest field of sources ever imaged in X-rays. The observation was performed with a wide field of view and therefore was best suitable for studying the overall demographics of the X-ray emission in the whole Orion Nebula Cluster. From the placement of the identified sources in the Hertzsprung-Russell (HR) diagram it was concluded that a large fraction of the cluster members with masses above 1 \\Msun are detected in X-rays, whereas the detection rate decreases rapidly for lower mass stars.  In our observations we focus on a smaller field of view around the Trapezium, which suffers from saturation effects in the Garmire et al.\\ (2000) data. Our objective is to resolve the Orion Trapezium cluster core in X-rays and to perform a detailed spectral analysis on the bright sources as well as to study the variability of the X-ray emission. We present the results from the zero order field analysis in three parts. Paper I gives an overview of these observations but focuses on the X-ray properties of the Orion Trapezium and its very close vicinity within a radius of 15\". Paper II (Kastner et al. 2000, in preparation) then discusses identifications within the whole available field of view around the Orion Trapezium with known infrared and radio catalogs as well as with newly obtained IR observations. In paper III (Schulz et al. 2000a, in preparation) we present X-ray spectroscopy on the whole field of view. Schulz et al. (2000) also present high resolution spectroscopy of $\\theta^1$ Ori C using the first order spectra of the HETGS. ",
        "conclusions": "Our two observations of the Orion Trapezium Cluster in 0th order of the Chandra HETGS with a spatial resolution of 0.5\" thoughout almost the entire field of view reveal many new details about the X-ray properties of the Cluster. They allow us to resolve the Trapezium Cluster entirely into point sources, study their variability patterns and perform detailed X-ray spectroscopy. Here we report on the following results:  We resolved all the emission around the Trapezium into point sources, detecting 111 X-ray sources. The cluster shows no diffuse emission above the total (both observations merged) sensitivity limit of 2.0$\\times10^{28}$ \\ergsec per arcsec$^2$. Besides the five main Trapezium stars we detected six additional X-ray point sources in the close vicinity of the Trapezium. Most sources are faint. Only 12 sources have persistent X-ray luminosities above $10^{31}$ \\ergsec; four of these are part of the Trapezium and are identified with early type stars. The avarage luminosity of the fainter X-ray sources is around 3$\\times 10^{30}$ \\ergsec, in agreement with the result from Garmire et al. (2000) for solar mass PMS stars.  The two observations showed considerable variability on timescales larger than 0.5 days. Fifteen percent of the sources vary by more than a factor 5 in luminosity; the largest excursion in luminosity was from 1.3$\\times 10^{31}$ \\ergsec to 4.1$\\times 10^{29}$ \\ergsec . Most of these variable sources appear absorbed with column densities between 1.7$\\times 10^{21}$ cm$^{-2}$ to 1.1$\\times 10^{24}$ cm$^{-2}$ with standard deviations of 5$\\%$ and 10$\\%$ respectively. The three bright Trapezium stars display very similar values of line of sight column density, with a mean value of $N_H =$ 1.93$\\pm0.29 \\times 10^{21} cm^{-2}$.  The most luminous source is $\\theta^1$ Ori C with L$_x$ =  1.85$\\pm0.18 \\times 10^{32}$ \\ergsec. This is an order of magnitude lower than the estimated luminosity from \\asca of $10^{33}$ \\ergsec. If we integrate the luminosities of all sources within the \\asca point spread function, however, we get a reasonably consistent result. Stahl et al. (1993) measured a period of 15.4 days in optical and UV data, which was also seen in the \\ros HRI data (Caillault et al. 1994). From the ephemeris given by Stahl et al. we determine phases of 0.68 for our first, and 0.23 for our second observation. From the HRI peak luminosity of given in Caillault et al. we expect luminosites of 2.3$\\times 10^{32}$ \\ergsec and 2.0$\\times 10^{32}$ \\ergsec, respectively, for our observations. Although the results in Table 1 are slightly lower than these estimates, they are not inconsistent given the various systematic uncertainties.  \\subsection{Early Type Stars}  Efforts to model X-ray spectra from stellar winds have been quite modest since the first observations with \\ein (Seward et al. 1979). The main reason is that the theory of line radiation-driven stellar winds, which explains successfully the mass loss from a hot luminous star from UV data, could never correctly predict observed X-ray fluxes. Therefore to date no definitive models for the production of X-rays in stellar winds exist. The X-ray emission is now thought to be produced by shocks forming from instabilities within a radiatively driven wind (Lucy $\\&$ White 1980), where hot blobs plow through rarified cool wind material well into the terminal flow of the wind (Lucy 1982). This model produces X-rays of about 0.5 keV temperature. Hellier et al. (1993) fitted high signal-to-noise PSPC spectra of $\\zeta$ Pup with detailed NLTE models under the assumption that the X-rays arise from shocks distributed throughout the wind and allowing for recombination in the outer regions of the stellar wind. The best fits predicted two temperatures of 1.6 and 5.0$\\times10^6$ K with shock velocities around 500 km s$^{-1}$.  We performed detailed spectroscopy on the Trapezium stars as well as for most of the nearby X-ray sources.  $\\theta^1$ Ori A and E are the brightest sources besides C, with persistent X-ray luminosities above $10^{31}$ \\ergsec. None of these sources are observed to be variable by more than 10$\\%$. All three sources are identified with massive, very early type stars: $\\theta^1 Ori$ C is a massive ($\\sim 25$ \\Msun) O7V star with a recently discovered low-mass companion (Weigelt et al. 1999); $\\theta^1 Ori$ A is identified with an eclipsing binary of 0.2\" separation and a 63.43 day period; $\\theta^1 Ori$ E is a massive early B0.5 type star, and so far no companion has been detected. The main component of $\\theta^1 Ori$ A is not very well determined and speculations range from spectral type B0V to B1V (Lloyd  $\\&$ Strickland 1999). The nature of the companion is similarly unclear and identifications range from a low-mass pre-main sequence star (Strickland $\\&$ Lloyd 2000) to a B8-AOV type star (Vitrichenko 1999). There are indications that $\\theta^1 Ori$ A is even a triple system (Petr et al. 1998). The spectra of $\\theta^1 Ori$ A, C, and E were fitted with a two-temperature Raymond-Smith spectrum of kT $\\sim$ 0.8 keV and kT $\\sim$ 2-3 keV. Since these sources are identified with very early type, massive stars we conclude that the X-ray spectra show X-ray emission characteristics of strong stellar winds. Specifically, in the case of $\\theta^1 Ori$ A it is more likely that the X-rays originate from its very early type main component and not from any companion. In fact if the companion is a T Tauri star, its X-ray spectrum may only contribute a small fraction of the observed flux. If it is a late type B-star, we would not expect any X-rays at all (Bergh\\\"ofer $\\&$ Schmitt 1994). The spectral fits also agree with the results from previous \\asca observations (Corcoran et al. 1994, Cohen et al. 1996) that X-rays from early type stars have two or more temperature components.  The sources $\\theta^1$ Ori B and D have luminosities well below $10^{31}$ \\ergsec. The fainter one, $\\theta^1$ Ori D, shows only a soft component of 0.7 keV and we may speculate that here the wind lacks the hot component observed in $\\theta^1$ Ori A, C, and  E, which are of earlier type. The \\ros spectrum of $\\lambda$Sco and $\\alpha$ Vir, which are of similar spectral type (B1V), show similar luminosities and spectral temperatures. Here we may observe an effect, that with increasing spectral type not only the X-ray luminosity decreases but also the spectrum softens. $\\theta^1$ Ori B is brighter, lacks the soft component, and shows a hot component. The second member of $\\theta^1$ Ori B, $\\theta^1$ Ori B1, whose stellar type is not precisely known,  seems to be the dominating source for the observed X-rays. Star B is identified with a spectral type of B3V (Abt et al. 1991), which may already be late enough to have such a weak wind that it cannot produce the X-ray luminosities observed. \\ros observations indicate that in B-stars of later type, the luminosity decreases. Cassinelli et al. (1994) analysed  several B-stars of this particular spectral type and found soft X-ray luminosities of the order of $10^{28}$ \\ergsec and lower, which is well below our sensitivity limit. On the other hand, star B1 is thought to be of earlier type (B1) and \\ros observations predict X-ray luminosities of above $10^{30}$ \\ergsec, which we do not observe. It also remains unclear why we observe such hard spectrum.  The spectra of the early type stars in the Trapezium seem unusually hot compared to what stellar wind models currently predict. Four of the five main components, A, B, C, and E, show spectral temperatures between 2.4$\\times10^7$ K and 3.7$\\times10^7$ K, which is an order of magnitude hotter than current stellar wind models predict. In the case of $\\theta^1$ Ori C, Schulz et al. (2000) recently determined a temperature range of 0.5 to 6.1$\\times10^7$ K from X-ray emission lines. These temperatures were deduced from H-like ion species up to S XVI and Fe XXV, which  rules out a significant inverse Compton continuum as suggested by Chen $\\&$ White (1991). Babel $\\&$ Montmerle (1997a) presented the hypothesis that in the case of $\\theta^1$ Ori C a dipolar magnetic field is embedded in the radiation-driven wind and X-rays are produced by a magnetically confined wind shock. In the case of the hot A0p star IQ Aur,  Babel $\\&$ Montmerle (1997b) modelled temperatures up to 9$\\times10^6$ K by including the effects of magnetic confinement. If in addition substantial stellar rotation was included, temperatures could be as high as 3 $\\times10^7$ K. Therefore it seems compelling to include magnetic field effects as well as centrifugal forces into stellar wind models for these hot stars. The fact that such high temperatures are present in $\\theta^1$ Ori A, B, and E, as well as in $\\tau$Sco (Cohen et al. 1997), also shows that these effects may be more common in early type stars.  \\subsection{YSOs}  Although it is now widely accepted that X-ray emission from young stars emerges from hot stellar coronae heated via a stellar dynamo, the underlying physics of this mechanism is still poorly understood. High levels of magnetic activity in connection with rotation, accretion, and/or outflows may play a key role (see Feigelson $\\&$ Montmerle 1999 for a review). Recent imaging observations showed that the X-rays of most T Tauri stars vary on time scales of days and a substantial number of stars display high amplitude flares with time scales of hours. X-ray spectra show  more than one temperature and are usually modelled with soft spectral components of 2-5$\\times10^6$ K and some hard components during flares of up to 3 $\\times10^7$ K. Montmerle et al. (2000) recently modelled the X-ray emission from a giant flare with a peak temperature of 7 $\\times10^7$ K in YLW 15 observed with \\asca (Tsuboi et al. 2000) with strong magnetic shearing and reconnection between the central star and the accretion disk. In the case of WL 6 (Kamata et al. (1997) the X-ray flaring is rotationally modulated (see also Stelzer et al. 1999). Montmerle et al. (2000) inferred a mass-rotation relation, in which higher mass stars are fast rotators and giant hard flares are generated through strong star-disk magnetic interactions, while lower mass stars are slower and and their X-ray emission is solely of stellar origin.  The vicinity of the Orion Trapezium harbours a large variety of YSOs. Most notable are the large number of proplyds detected with the Hubble Space Telescope (O'Dell et al. 1993). We detected 8 faint X-ray sources, where six sources are identified with radio objects, five sources with proplyds in the vicinity of $\\theta^1$ Ori C (i.e. proplyds 163-317, 166-316, 167-317, 168-328, and 171-334). One source is identified with a YSO in the the IR only. These X-ray sources show luminosity variations between a factor 3 and 10, and have spectra indicating that they are heavily absorbed. The X-ray emission, together with optical, IR and radio surveys (Bally et al. 1998, McCaughrean and Stauffer 1994, Felli et al. 1993) indicate that these sources are embedded Class I, Class II,  or Class III protostars. We do not believe that the X-rays are generated through induced shocks by the wind $\\theta^1$ Ori C. First, the flux of these sources varies by large factors which we would not expect under constant illumination unless the density structure in the wind shock at the proplyds changes similarly. That would imply similar brightness variations in the optical, but no such variations have been reported so far. On the other hand, in source $\\#$78 we see variability timescales of the order of 10 min to 2 hours (figure 2). From this we deduce an upper limit of the  size of the emitting region of the order of 1.0 - 12 AU, which is too small compared to size of the proplyds ($\\sim$  500 AU, Bally et al. 1998). Second, the spectra we analysed are intrinsically  absorbed. Thus the emission likely originates from deep inside the dusty disks, rather than from the wind shock region.  In the context of recent observations of hard X-ray emission from YSOs, these results are quite remarkable; the light curves indicate that only one of the sources actually displays large flaring activity, and the others are persistent in flux with some minor activity. Very large flux variations occur only on time scales longer than 12 h. The spectra, however, indicate extremely high temperatures at quite modest luminosities, from which we conclude that the X-ray emission may not be related to star-disk magnetic interactions in connection with rapid rotation as proposed by Montmerle et al. (2000). The observed X-ray behavior of these proplyds seem more related to recent \\asca observations of SU Aurigae (Skinner $\\&$ Walter 1998), which has been identified with a type G2  CTTS and where it was suggested that the X-ray emission arises from a magnetically confined stellar plasma. In this respect long term X-ray observations of these proplyds will not only enhance the significance of these hot spectra but also  allow a search for rotational modulations and lead to better understanding of the magnetic properties and their role in the evolution of these protostars."
    },
    "0011/astro-ph0011372_arXiv.txt": {
        "abstract": "We have observed the canonical gigahertz-peaked spectrum source 0108+388 with the VLBA at a range of frequencies above and below the spectral peak. The activity that dominates the radio emission from 0108+388, which is also classified as a Compact Symmetric Object, is thought to be less than 1000 years old.  We present strong evidence that the spectral turnover in 0108+388 results from free-free absorption by non-uniform gas, possibly in the form of a disk in the central tens of parsecs. ",
        "introduction": "The radio source 0108+388 (RA=$01^{\\rm h} 11^{\\rm m} 37^{\\rm s}.32$, Dec=$39^\\circ 06'28''.1$ (J2000), $z$=0.699) is an example of a class of extragalactic radio source referred to as Gigahertz-Peaked Spectrum (or GPS) sources.  GPS sources, which are discussed at length in a thorough review by O'Dea (1998), are characterized by large radio luminosities, compact structure, and spectra that turn over at Gigahertz frequencies with steep spectra on either side of the peak (O'Dea, Baum, \\& Stanghellini 1991; O'Dea 1998).  The steep slope at high frequencies is generally accepted to be due to synchrotron losses, but there are two main competing theories for the cause of the low frequency turnover.  Either the lower frequency radio emission is absorbed by ionized gas in the line of sight in the host galaxies, with conditions comparable to that of emission-line regions in active galaxies (van Breugel, Heckman, \\& Miley 1984, Bicknell, Dopita, \\& O'Dea 1997; Kameno et al.\\ 2000), or the nuclei of these galaxies have exceptionally large magnetic fields, causing synchrotron self-absorption at GHz frequencies (Hodges, Mutel, \\& Phillips 1984; Mutel, Hodges, \\& Phillips 1985; O'Dea et al.\\ 1991; Readhead et al.\\ 1996; de Vries, Barthel, \\& O'Dea 1997; Snellen et al.\\ 2000). Absorption by induced Compton scattering in combination with free-free absorption has also been proposed by Kuncic, Bicknell, \\& Dopita (1998).  As GPS sources have been studied in more detail since the discovery of a large population of such sources in the 1980's (Gopal-Krishna, Patnaik, \\& Steppe 1983, Spoelstra, Patnaik, \\& Gopal-Krishna 1985), it has become apparent that they constitute a diverse group, with morphological and spectral variances.  The object we discuss here, 0108+388, is a canonical case of a morphological group referred to by Wilkinson et al.\\ (1994) as compact symmetric objects (or CSO's), which have small ($<$ 1 kpc) symmetric radio emission structure. Multi-epoch VLBI observations of a number of CSO's have revealed expansion rates that imply ages of the order of 1000 years (Owsianik, Conway, \\& Polatidis 1998; Owsianik \\& Conway 1998; Taylor et al.\\ 2000), supporting the earlier hypothesis that CSO's are precursors to the young classical double radio galaxies (Phillips \\& Mutel 1982; Carvalho 1985; Readhead et al.\\ 1996).  VLBI observations of 0108+388 by Owsianik et al.\\ (1998) and by Taylor et al.\\ (2000) infer apparent separation speeds of the outer components of 0.197 $\\pm$ 0.026 $h^{-1}$ c and 0.24 $\\pm$ 0.04 $h^{-1}$ c, respectively, from which these authors derive kinematic ages of just 367 $\\pm$ 48 yr and 310 $\\pm$ 70 yr, respectively.  This is the youngest age yet measured for a radio galaxy.  However, with VLA and WSRT observations of 0108+388 Baum et al.\\ (1990) detected extended emission 20 arcsec (78 $h^{-1}$ kpc) from the core, casting an apparent contradiction to the young classical double radio galaxy model.  Baum et al.\\ propose that either the activity in 0108+388 is recurrent or has recently been smothered by infalling gas.  High resolution imaging of 0108+388 was first undertaken at 5 GHz as part of the PR VLBI survey (Pearson \\& Readhead 1988) and found to have a simple double structure.  Conway et al.\\ (1994) and Taylor, Readhead, \\& Pearson (1996) made more sensitive and higher resolution images at higher frequencies that revealed a steeply inverted core component connected to both outer components by a faint chain of components.  Overall, 0108+388 has an `S' symmetry that is relatively common among CSOs (Readhead et al.\\ 1996).  The spectrum of the total radio emission from the compact structure in 0108+388 was analyzed by Baum et al.\\ (1990).  They measured a spectral index below the turnover that was steep and roughly constant, in conflict with the exponential turnover expected from free-free absorption.  However, their observed spectrum is the total from several components in the source.  An exponential turnover in the free-free absorption case will apply only if the absorbing medium uniformly covers the entire radio source.  We report here on a set of VLBA observations designed to study the absorption at each point in the map.  We have observed 0108+388 with the VLBA at five frequencies: two above the turnover, two below the turnover, and one at the turnover frequency itself.  The maps made with these observations reveal the morphology of the radio emission from 0108+388 and how it is obscured by absorption.  We find that the spectra of 0108+388 at all positions in our maps are consistent with free-free absorption. Additionally, our maps reveal that the absorption is non-uniform and is suggestive of an edge-on disk centered on the core.  Other evidence for a large amount of thermal gas along the line of sight to the core of 0108+388 has been reported by Carilli et al.\\ (1998), who measured significant absorption by H{\\sc i}, and by Cawthorne et al.\\ (1993), who find no measurable polarization of the radio emission. The existence of a circumnuclear gas disk, with likely instabilities, could also be the source of intermittent fueling of the central engine and explain the recent onset of activity along with the past activity. ",
        "conclusions": "By extrapolating our high frequency maps, we obtained opacity maps of 0108+388 at two frequencies below the spectral turnover which strongly suggest that free-free absorption, possibly by a disk of ionized gas, is indeed the cause of the spectral turnover, and perhaps in other CSOs with a GHz-peaked spectrum.  The spectra all across the map fit free-free absorption well.  Reasonable values for temperature (10$^4$ K) and electron density (500 cm$^{-3}$) in a disk of order 100 pc in size fit the observed optical depths.  The spectrum of the core, though, is inverted even at our highest frequencies, probably because of synchrotron self-absorption.  The presence of a dense gas disk in 0108+388, which is also consistent with the large column densities implied by the H{\\sc i} absorption measurements of Carilli et al.\\ (1998), may be related to the recent onset of activity in this object, as indicated by the age estimates by Owsianik et al.\\ (1998) and Taylor et al.\\ (2000).  The accumulated evidence thus strongly supports the conclusions that CSOs in general are intrinsically young objects that evolve into classical kpc-scale radio galaxies, and that the radio galaxy 0108+388, in particular, is the youngest source known."
    },
    "0011/astro-ph0011144_arXiv.txt": {
        "abstract": "s{  \\begin{center} {\\bf Abstract} \\end{center} The interaction between the  Earth's magnetic field and the solar wind plasma results in a natural plasma confinement system which stores energy. Dissipation of this energy through Joule heating in the ionosphere can be studied via the Auroral Electrojet (AE) index. The apparent broken power law form of the frequency spectrum of this index  has motivated investigation of whether it can be described as fractal coloured noise.  One frequently-applied test for self-affinity is to demonstrate linear scaling of the logarithm of the structure function of a time series with the logarithm of the dilation factor $\\lambda$. We point out that, while this is conclusive when applied to signals that are self-affine over many decades in $\\lambda$, such as  Brownian motion, the slope deviates from exact linearity and the conclusions become ambiguous when the test is used over shorter  ranges of $\\lambda$. We demonstrate that non self-affine time series made up of random pulses can show near-linear scaling over a finite dynamic range such that they could be misinterpreted as being self-affine. In particular we show that pulses with functional forms such as those identified by Weimer within the $AL$ index, from which $AE$ is partly derived, will exhibit nearly linear scaling over ranges similar to those previously shown for $AE$ and $AL$. The value of the slope, related to the Hurst exponent for a self-affine fractal, seems to be a more robust discriminator for fractality, if other information is available. } ",
        "introduction": "The characterisation of global energy transport in the coupled solar wind-magnetosphere-ionosphere system is a fundamental problem in  space plasma physics~\\cite{ke}. Solar wind energy is transferred to, stored by, and ultimately released from the magnetosphere by a range of mechanisms, in which substorms play a key role. Most investigations of the substorm problem have focused on single substorms or small groups of similar events, analogous to the study of individual earthquakes in seismology.  A complementary approach is to analyse inputs to and outputs from the system in an attempt to constrain the range of possible physics occurring in the magnetospheric ``black box'' (c.f. analogous approaches in climatology and seismology~\\cite{tu}). Reviews of the significant progress made so far in applying the methods of low dimensional chaos to the magnetosphere are given by Klimas {\\it et al}~\\cite{kl} and Sharma~\\cite{sh}; while more recent investigations into whether or not the ``black box'' can be treated as a self-organised critical (SOC) system~\\cite{je} are reviewed by Watkins {\\it % et al}~\\cite{wat00}, Chapman and Watkins~\\cite{chap00} and Consolini and Chang~\\cite{con00}. One mechanism for dissipation of magnetospheric energy is through Joule heating in the ionosphere's auroral electrojets. This process can be studied via the auroral electrojet ($AE$) index, which is a means of estimating the electrojet current. The Joule energy dissipated depends upon both this and the ionospheric conductivity. $AE$ is available at 1-minute resolution. Tsurutani {\\it et al}.~\\cite{ts} showed this to have a ``broken power law'' frequency spectrum. The high frequencies approximately follow $f^{-2}$ while the lower frequencies are $f^{-1}$ with a break at about $1/5$ h$^{-1}$. Power law frequency spectra are common in nature and can have several causes~\\cite{je} such as Kolmogorov turbulence or the bifurcation route to chaos. They are thus in themselves not sufficient to completely constrain simple models. A parallel effort to studies of the power spectrum has been the search for low dimensionality, initially through the Grassberger-Procaccia (GP) algorithm~\\cite{kl,sh}. However, as noted by Osborne and Provenzale~\\cite{op}, a low and fractional GP dimension is not uniquely a signature of low dimensional chaos. It is also compatible with self-affine coloured noise~\\cite{op} or SOC~\\cite{ch}. In view of the fact that $AE$ is known {\\it a priori} to be the output of a complex system, Takalo and Timonen, in an important series of papers$^{12-16} $, % investigated whether the dynamics of magnetospheric and auroral indices were better encapsulated by stochastic ``coloured noise'' rather than by chaos. One test applied to $AE$~\\cite{tt94} was for self affinity - a property of both coloured noise and chaos. A particularly important technique for identifying self-affinity in the work of Takalo and Timonen$^{12-16}$ was the use of the second order structure function $S_{2}$ (although other methods have also been applied to this problem$^{17-19}$). In this paper, by constructing a simple example, we illustrate that $S_{2}$ alone cannot reliably distinguish exponential autocorrelation from intrinsic self-affinity in the short timescale part of the $AE$ signal, which has been linked to the substorm ``unloading'' timescale~\\cite{tt98}. By considering how $S_{2}$ is related to other measures of self-affinity we address the question of what additional knowledge may be required to make $S_{2}$ more useful.  \\begin{figure}[t] \\psfig{figure=fig1.eps,height=3.0in} \\caption{ Power spectrum of model Brownian motion ($H=0.5$). } \\label{fig:radish} \\end{figure} ",
        "conclusions": "We now consider the scaling properties of the measured $AE$ time series in the light of the previous examples. The top panel of figure 6 shows the autocovariance of the first 100 days of $AE$ for 1983, and may be compared with the top panel of figure 4 of Takalo and Timonen~\\cite{tt94}. Again, the steep (exponential) fall of the ACV results in a near linear slope for small $\\lambda \\Delta t$, and a slow decrease in the slope as larger and larger ranges of $\\lambda \\Delta t$ is considered . Importantly, however, the slope is, always less than 1 (J. Takalo, Private communication, 1999). Overall it resembles near-linear scaling in the structure function for most of the first two decades of $\\lambda$ (bottom panel of figure 6), and (plotted in the the middle panel of figure 4 of Takalo and Timonen~\\cite{tt94}), was cited by Takalo and Timonen~\\cite{tt94} as a key piece of evidence for self-affinity in $AE$. They also noted the resemblance of the $AE$ autocovariance function to an exponential and proposed that the autocorrelation time of $AE$ be defined as the lag for timescales longer than that over  which the autocovariance ceased to be exponential. Inspection of figures 4, 5 and 6 lead us to conclude, however, that, unlike the ideal case of Brownian motion, {\\bf neither} the curve of $S_2$ for $AE$% , {\\bf nor} that of the simplified random Weimer pulse train are straight over the range $\\lambda \\Delta t=1$ to $\\lambda=120$. We remark that, insofar as the ACF of $AE$ is exponential for small $\\lambda \\Delta t$, there {\\em must} eventually be a departure from near-linearity in the structure function as $\\lambda \\Delta t$ increases, unless the range over which the exponential behaviour is seen is so small that a straight line would be just as good an approximation as the exponential.  In addition, both $AE$ and the model Weimer pulse train of section 4.2 give a fractional $H$ value when taken over the whole range from $10$ to $100$. Without {\\em a priori} additional knowledge, we might equally well have concluded that the random pulse train was self-affine over the range $% \\lambda \\Delta t < 100$, but by construction we know this is not so.  Our model was deliberately simplified. In the natural $AE$ time series, the extended tail of the ACV is  expected to reflect the solar wind-driven component (also present in $AU$ and $AL$), which our simulation neglected. As originally conjectured by Tsurutani {\\it et al}~\\cite{ts} the solar wind driver is probably the origin of the ``1/$f$\" part of the $AU/AL/AE$ spectrum~\\cite{fre00,tak00}.  We may summarise our findings as follows. By construction of an explicit counter-example we have shown that near-linear scaling of $S_2$ over about two decades is not in itself sufficient to show self-affinity. We have also given analytic and numerical evidence that non-fractal random series can produce non-integer Hurst exponents over limited dynamic ranges. We thus infer that self-affinity in the range $0$ to $100$ minutes for $AE$ has not been and could not be proved by the use of $S_2$ alone.  One may reasonably point out that several other methods have been used to provide evidence of self-affinity in geomagnetic indices; both in the papers of Takalo et al. and those of other workers$^{17-19}$. One may thus enquire as to what kind of additional knowledge or analysis techniques would be necessary for considering the results of the structure function method fruitful? Based on what we have found, we suggest that answer is at least threefold.  1) {\\em Be aware that many tests for fractality are actually designed assuming a fractal signal}: A test based on the assumption of fractality can disprove fractality but cannot prove it. The methods for measuring fractal dimension that we are aware of assume self-affinity in their design  i.e. they typically examine the scaling behaviour of a signal. Only if they find no evidence of scaling at  all is there no ambiguity.  2) {\\em Use more than one test}: Several tests are better than one because different methods are sensitive to different  non-fractal effects. Thus use of several tests means that a series  with non-fractal aspects is less likely to be misinterpreted.  Most of the methods for measuring fractal dimension which have been applied to geomagnetic data are of one of two basic types. The first type of method basically estimates the dimension of a fractal curve by examining how the average value of short lengths of curve \\begin{equation} S_1 = < X(t + \\lambda \\Delta t) - X(t)> \\end{equation} scales with the ruler length $\\lambda$ (in units of the sampling interval $% \\Delta t$). Such methods have been applied by V\\\"{o}r\\\"{o}s~\\cite{vor90} to magnetometer data, and more recently to geomagnetic and solar wind quantities by Price and Newman~\\cite{price01}, who used the related, cumulative \"R/S\" analysis.  The second type studies the positive definite second order function \\begin{equation} S_2 = <(X(t + \\lambda \\Delta t) - X(t))^2> \\end{equation} and returns the same information~\\cite{ro} as the ACF when estimated on a stationary signal (see section 2).  For this reason it is thus also formally related to the power spectrum via  the Wiener-Khinchine theorem. $S_2$, the ACF and the power spectrum have all been extensively investigated for the $AE $, $DSt$ and related indices by Takalo {\\em et al}$^{12-16}$ .  The meaning of this family of techniques can be understood as studying the behaviour of the histogram of variance of the signal (or the power spectrum) with increasing time dilation $\\lambda \\Delta t$ (or frequency); depending on whether one is dilating in time (in the case of $S_2$ and the ACF) or frequency (in the case of the power spectrum). We caution that time lag in the ACF or in $S_2$ is not trivially 1/(the Fourier frequency) because any frequency in a Fourier transform has contributions from multiple ACF lags and vice versa (see Bendat and Piersol~\\cite{be86}, pages 120-122). In the case of a simple fractal, the dimension (and Hurst exponent $H$) estimated from such methods should theoretically be the same as from $S_1$, although in practice the errors of the two methods need not be the same~\\cite{san92}. If they differ substantially, this  may be a pointer that the time series is not intrinsically a wideband fractal, and that one of $S_1$ or $S_2$ is more sensitive to this.  An example of how additional tests for fractality have supplied new knowledge is in the continuing  study of the $AE$ index. This has been known since the work of Tsurutani {\\em et al}~\\cite{ts} to have a  ``1/$f$\" low frequency and ``1/$f^2$\" high frequency power spectral density. Acting only on information from the power spectrum or other $S_2$-type methods, one might thus infer that AE is a bi-affine quantity$^{12-16}$, i.e. it has two separate scaling regions and a break between them. In contrast, Consolini~\\cite{con97} has studied the ``burst distributions\"~\\cite{con99,con97} of $AE$. These are the histograms of intervals between threshold crossing times (burst and inter-burst lifetimes) and of areas above threshold between crossings (burst sizes). The lifetime distributions are an $S_1$-type measurement~\\cite{addison,fre00b} and were found to have (exponentially rolled-off) scaling with a single slope over a very wide range, interrupted only by a non-scaling component at about 2 hours. The apparently paradoxical observation of bi-affine behaviour in $S_2$ and ``contaminated\" mono-affine behaviour in $S_1$ has been addressed in two different ways. One has been to introduce models which have the required properties in both $S_1$ and $S_2$, such as forest fire models~\\cite{con99} or coupled map lattices~\\cite{tak00}  driven by wideband solar-wind like signals. The other, informed by the fact (section 4 and 5 above) that the high frequency $f^{-2}$ part of a power spectrum {\\em need not arise from a fractal aspect of the time series}, has been to postulate~\\cite{fre00} that the $AE$ series is in fact a hybrid time series with a fractal element arising from the solar-wind driven ionospheric current systems and a non-fractal part arising from energy storage and release in the magnetosphere (substorms). This was supported~\\cite{fre00} by the observation that the scaling in $AE$ (and $AU/AL$) burst lifetimes is the same as that seen in the solar wind (see also Freeman et al~\\cite{fre00b}) while the non-scaling component was seen only in the magnetospheric quantities such as  $AU$ and $AL$~\\cite{fre00}.  There have been exceptions to the use of $S_2$ or $S_1$ type techniques in the geomagnetic context. We are grateful to an anonymous referee for reminding us of the results of a multifractal analysis of the $AE$ index by Consolini et al.~\\cite{con96}. These results must imply some constraints on possible models describing the variability of auroral currents. However, in the same way that $AE$ when measured over 1 year by a method of type $S_1$ is essentially fractal~\\cite{con97}, and required the use of several years' measurements for the ``bump\"-like feature in the otherwise scaling $S_1$ to become apparent~\\cite{fre00}, it seems to us that one might expect a multifractal  analysis of less than 2 months of $AE$~\\cite{con96} to give a good fit to a p-model of turbulence because the solar wind driver is also well fitted by this particular turbulence model~\\cite{horbury}. We believe that a  study on a much longer series of $AE$ would be required to exclude even our own toy model of random differentiable Weimer pulse trains, when superposed on the multifractal solar wind background. We note that, independently, a recent multifractal study of geomagnetic data from Thule, Alaska has excluded the biaffine coloured noise model~\\cite{voros00} for that dataset.  3) {\\em Remember that Nature does not have to be purely fractal any more than it has to be non-fractal}: Many types of natural signal have both fractal {\\bf and} non-fractal components. In consequence, when using methods to examine fractality, one should be aware that it is possible to find something between the extremes of wideband fractality and none at all, as  discussed in point (2) above  for the case of $AE$. Another example is to imagine looking out of one's tea-room window at a tree through a regularly spaced window blind. The distinguishing of the fractal tree and the periodic blind is a task that the human eye and brain perform routinely, and which a Fourier transform can also do because it can resolve the blind spacing as a spatial frequency. A ``random blind\" appearing at Poisson-switched intervals would be much more of a problem for an FFT, and would be analogous to the pulses of section 3 and 4. The user thus needs to determine how much the presence of a ``contaminating\" signal or signals in the fractal time series may affect their interpretation, at which point the question may become as much physical as mathematical. This is currently an admittedly very difficult task because of the sparsity  of literature on such hybrid time series, and is one which we plan to examine in more detail in future papers."
    },
    "0011/astro-ph0011234_arXiv.txt": {
        "abstract": "We investigate the Mg--$\\sigma$ and \\fe\\/--$\\sigma$ relations in a sample of 72 early-type galaxies drawn mostly from cluster and group environments using a homogeneous data-set which is well-calibrated onto the Lick/IDS system. The small intrinsic scatter in Mg at a given $\\sigma$ gives upper limits on the spread in age and metallicity of $49$\\% and $32$\\% respectively, if the spread is attributed to one quantity only and if the variations in age and metallicity are uncorrelated. The age/metallicity distribution as inferred from the \\Hb\\/ {\\em vs}\\/ \\fe\\/ diagnostic diagram reinforces this conclusion, as we find mostly galaxies with large luminosity weighted ages spanning a range in metallicity.  Using Monte-Carlo simulations, we show that the galaxy distribution in the \\Hb\\/ {\\em vs}\\/ \\fe\\/ plane cannot be reproduced by a model in which galaxy age is the only parameter driving the index-$\\sigma$ relation. In our sample we do not find significant evidence for an anti-correlation of ages and metallicities which would keep the index--$\\sigma$ relations tight while hiding a large spread in age and metallicity. As a result of correlated errors in the age-metallicity plane, a mild age-metallicity anti-correlation cannot be completely ruled out given the current data. Correcting the line-strengths indices for non-solar abundance ratios following the recent paper by Trager~et~al., leads to higher mean metallicity and slightly younger age estimates while preserving the metallicity sequence. The [Mg/Fe] ratio is mildly correlated with the central velocity dispersion and ranges from [Mg/Fe]~$=$~0.05 to 0.3 for galaxies with $\\sigma > 100$~\\kms. Under the assumption that there is no age gradient along the index--$\\sigma$ relations, the abundance-ratio corrected Mg--$\\sigma$, Fe--$\\sigma$ and \\Hb--$\\sigma$ relations give consistent estimates of $\\Delta {\\rmn [M/H]}/ \\Delta \\log \\sigma \\simeq 0.9 \\pm 0.1$. The slope of the \\Hb--$\\sigma$ relation limits a potential age trend as a function of $\\sigma$ to 2-3 Gyrs along the sequence. ",
        "introduction": "\\label{sec:intro} Over the last two decades spectroscopic and photometric observations of nearby early-type galaxies have shown that they obey tight scaling relations. For example more luminous galaxies are redder \\cite{san78,lar80,bow92}; the strength of the Mg-absorption (at $\\sim$5175 \\AA) increases with increasing central velocity dispersion \\cite{ter81,ben93,coll99}; and early-type galaxies populate only a band (or `plane') in the three dimensional space of central velocity dispersion, effective radius, and mean effective surface brightness, the so called `Fundamental Plane' \\cite{djo87,dre87}.  The tightness of the scaling relations has traditionally been interpreted as evidence for a very homogeneous population of early-type galaxies, \\ie\\/ old stellar populations and similar dynamical make-up of the galaxies \\cite{bow92,ren93,ell97,vdo98}. For example the colour-magnitude relation is perhaps best explained by a correlation of the metal abundance of the stellar population and total galaxy mass with no age gradient along the sequence \\cite{bow92,kod98,vaz2000}. However, evidence is being found that elliptical galaxies do generally have complicated dynamical structures (\\eg\\/ decoupled cores) and disturbed morphologies (\\eg\\/ shells), suggesting an extended and complex assembly such as in a hierarchical merging picture (\\eg\\/ Cole et~al. 2000). Furthermore measurements of absorption line-strength indices together with evolutionary population synthesis models suggest that the mean age of the stellar populations in early-type galaxies span a wide range \\cite{gon93,tra98}.  There is a hint of a connection between the detailed morphology and the mean ages of the stellar populations as ellipticals with disky isophotes show on average younger ages \\cite{dej97,tra97}. This is in agreement with the studies of cluster galaxies at modest look-back times of 3-5 Gyr where an increased fraction of blue, late-type galaxies is found \\cite{but78,but84}. Probably this change is associated with the transformation of spiral galaxies into S0 types observed over the same interval \\cite{dre97,cou98}. How can we reconcile the existence of tight scaling relations for present day early-type galaxies with the evidence of an extended galaxy assembly and star formation history? An answer to this question rests largely on our understanding of the detailed physical processes which determine the spectral or photometric properties seen in scaling relations.  Broad band colours and line-strength indices of integrated stellar populations are sensitive to changes in both their age and metallicity. In fact it is possible to find many different combinations of ages and metallicities which produce similar overall photometric properties (age-metallicity degeneracy, Worthey 1994). In order to explain tight scaling relations, despite the presence of a substantial diversity in the integrated stellar populations, it has been suggested that the ages and metallicities of integrated stellar populations `conspire' such that deviations from a given scaling relation due to age variations are balanced by an appropriate change in metallicity and vice versa; \\eg\\/ an age-metallicity anti-correlation \\cite{tra98,fer99,tra2000b}. For example, a re-analysis of the \\gon\\/ sample by Trager et~al. appears to show that galaxies with a young stellar component are also more metal rich. Recently, evidence for a negative age-metallicity correlation has also been found in other data sets \\cite{jor99} and in literature compilations \\cite{ter2000}. In contrast to this Kuntschner \\& Davies (1998) and Kuntschner (2000) did not find strong evidence for the existence of an age-metallicity anti-correlation in their analysis of the Fornax cluster.  However, to date there is a lack of large, high quality samples with which it would be possible to probe the population of early-type galaxies as a whole.  In this paper we seek to further our understanding of scaling relations such as Mg--$\\sigma$. We investigate the spread in the mean ages and metallicities {\\em at a given mass}\\/ but also probe the possibility of age and metallicity trends {\\em along}\\/ the relations. We use established methods such as the analysis of the scatter \\cite{coll99} in the Mg--$\\sigma$ relation and line-strength age/metallicity diagnostic diagrams, and take advantage of recent advances in stellar population modeling such as improved corrections for non-solar abundance ratios \\cite{tra2000a}.  We employ a homogeneous and high-quality subset of the data collected initially for the SMAC peculiar motion survey \\cite{hud99,smi2000}. The sample is not complete but includes galaxies drawn mostly from nearby clusters and groups with some galaxies from less dense environments.  The present paper is organized as follows. Section~\\ref{sec:data} describes the sample selection, observations, data reduction and the measurements of the line-strength indices. Section~\\ref{sec:mg_sig} presents the Mg--$\\sigma$ and Fe--$\\sigma$ relations with an initial investigation into the scatter and slopes. In Section~\\ref{sec:age} we present a detailed analysis of age-metallicity diagnostic diagrams including an improved treatment of non-solar abundance ratios. The effects of the non-solar abundance ratio corrections on the age-metallicity estimates and on the index--$\\sigma$ relations are presented in Sections~\\ref{sec:imp_index} and \\ref{sec:imp_sig} respectively. Our conclusions are given in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  For our sample of 72 early-type galaxies, drawn mostly from clusters and groups, we have analysed the \\mgbp\\/--$\\sigma$ and \\fep\\/--$\\sigma$ relations as well as age/metallicity diagnostic diagrams with up-to-date stellar population models. Taking the effects of non-solar abundance ratios into account, we conclude:  \\begin{itemize}  \\item Tight index-$\\sigma$ relations, as well as the results from age/metallicity diagnostic diagrams, provide evidence that cluster early-type galaxies are made of old stellar populations (age $>$7~Gyrs) and follow mainly a trend of increasing metallicity with increasing central velocity dispersion $\\sigma$. The galaxies display no or at most a mild age-trend along the sequence. A small number ($\\sim$5) of galaxies show younger luminosity weighted ages and above average metallicities in their mean stellar populations.  These galaxies have velocity dispersions below $\\sim$170~\\kms.  \\item Correcting the line-strength indices for non-solar abundance ratios is a crucial step forward in deriving more reliable age-metallicity estimates. Correlated errors, however, in age/metallicity diagnostic diagrams combined with the precision of currently available index measurements for medium to large samples lead to artificial anti-correlations between age and metallicity estimates. Therefore it is difficult to differentiate between mild anti-correlations between age and metallicity and a pure metallicity sequence. For the majority of our sample we do not find evidence for a prominent age/metallicity correlation and we can firmly exclude a pure age sequence at constant metallicity.  \\item Correcting the \\mgb\\/ and \\fe\\/ indices for non-solar abundance ratios changes the slope of the index--$\\sigma$ relations in the sense that the Mg--$\\sigma$ relation becomes shallower and the \\fe\\/--$\\sigma$ relation becomes steeper. The \\Hb--$\\sigma$ relation remains virtually unchanged. Under the assumption of no age gradients along the index--$\\sigma$ relations all three index--$\\sigma$ relations give consistent estimates of the total metallicity change per dex in $\\log \\sigma$: $\\Delta {\\rmn [M/H]}/ \\Delta \\log \\sigma = 1.00\\pm0.11$, $0.83\\pm0.11$ and $1.00\\pm 0.26$ for the \\mgb--$\\sigma$, $<$Fe$>$--$\\sigma$, and \\Hb--$\\sigma$ relations respectively. Introducing an age gradient along the sequence leads readily to inconsistent results of model predictions and the slope of the \\Hb--$\\sigma$ relation. At most an age gradient of 2 to 3 Gyr along the scaling relations is within our measurement errors.   \\item The [Mg/Fe] ratio is mildly correlated with the central velocity dispersion and ranges from ${\\rmn [Mg/Fe]} = 0.05$ to 0.3 for galaxies with $\\sigma \\ge 100$~\\kms. Some low luminosity galaxies, such as NGC~4464, do not follow the main trend but show [Mg/Fe] ratios similar to the more luminous galaxies.  \\end{itemize}"
    },
    "0011/astro-ph0011002_arXiv.txt": {
        "abstract": "We summarize recent developments in the study of the origin of halo spin profiles and preliminary implications on disk formation. The specific angular-momentum distributions within halos in N-body simulations match a universal shape, $M(<\\!j)\\!\\propto\\!j/(j_0\\!+\\!j)$. It is characterized by a power law over most of the mass, and one shape parameter in addition to the spin parameter $\\lambda$. The angular momentum tends to be aligned throughout the halo and of cylindrical symmetry. Even if angular momentum is conserved during baryonic infall, the resultant disk density profile is predicted to deviate from exponential, with a denser core and an extended tail. A slightly corrected version of the scaling relation due to linear tidal-torque theory is used to explain the origin of a typical power-law profile in shells, $j(M) \\propto M^s$ with $s\\gsim 1$. While linear theory crudely predicts the amplitudes of halo spins, it is not a good predictor of their directions. Independently, mergers of halos are found to produce a similar profile due to $j$ transfer from the orbit to the product halo via dynamical friction and tidal stripping. The halo spin is correlated with having a recent major merger, though this correlation is weakened by mass loss. These two effects are due to a correlation between the spins of neighboring halos and their orbit, leading to prograde mergers. ",
        "introduction": "Disk galaxies that form in cosmological hydro simulations are surprisingly smaller and of lower spin than observed disks, posing an ``angular-momentum problem\" (e.g. Navarro \\& Steinmetz 1997, 2000). A key ingredient in the origin of galactic rotation, which can be studied in detail, is the distribution of angular momentum in the dark halo within which baryonic infall is assumed to form the disk. We describe here several advances in the investigation of the origin of halo angular momentum and its distribution, and preliminary implications on disk formation, based on high-resolution $N$-body simulations and analytic approximations. The universal $j$ profile (\\S2) and the implied disk structure (\\S3) are presented in Bullock et al.~(2000, BDK+). The issues of tidal-torque theory (\\S4) are studied by Porciani, Dekel, \\& Hoffman (2000, PDH) and the scaling relation is by Porciani \\& Dekel (2000, PD). The $j$ profile due to mergers (\\S5) is by Dekel \\& Burkert (2000, DB). The correlations with merger history (\\S6) are to be described in Wechsler et al.~(2000, WBD+). ",
        "conclusions": ""
    },
    "0011/hep-ph0011251_arXiv.txt": {
        "abstract": "{\\rm We study how well KamLAND, the first terrestrial neutrino experiment capable of addressing the solar neutrino problem, will perform in ascertaining whether or not the large mixing angle MSW solution (with \\mbox{$10^{-5} \\lsim \\Delta m_{21}^2 \\lsim 10^{-4}\\, \\rm{eV}^2$} and oscillation amplitude $\\rm{sin}^2 2\\,\\theta_{12}>0.3$) is correct.  We find that in a year of operation KamLAND will provide unequivocal evidence for or against this solution. Furthermore, its sensitivity to the three-neutrino oscillation parameters in this region is sufficiently acute as to determine $\\Delta m_{21}^2$ to approximately $\\pm 10$\\% (for $\\rm{sin}^2 2\\,\\theta_{12}>0.7$)  and to fix $\\rm{sin}^2 2\\,\\theta_{12}$ to within $\\pm 0.1$ (at the $2\\sigma$ level) with three years of accumulated data, independent of the value of $\\theta_{13}$. } ",
        "introduction": "Several neutrino oscillation experiments now indicate that neutrinos are massive and that neutrino flavor mixing occurs. Atmospheric neutrino experiments (Kamiokande~\\cite{k1}, SuperKamiokande~\\cite{sk1}, IMB~\\cite{imb}, Soudan~\\cite{soudan} and MACRO~\\cite{macro}) report a $\\nu_{\\mu}/\\nu_{e}$ event ratio that is about 0.6 times the expected ratio; the $\\nu_{\\mu}$ flux shows a zenith angle dependence that is explained by oscillations. All experiments that measure the solar neutrino flux (Homestake~\\cite{homestake}, SAGE~\\cite{sage}, GALLEX~\\cite{gallex}, Kamiokande~\\cite{k} and SuperKamiokande~\\cite{sk}) find a deficit of $1/2$ to $1/3$ of the Standard Solar Model prediction \\cite{SSM}.  LSND~\\cite{lsnd}, an accelerator experiment, finds evidence for $\\bar{\\nu}_{\\mu}\\rightarrow \\bar{\\nu}_e$ and ${\\nu}_{\\mu}\\rightarrow {\\nu}_e$ oscillations, a result that is not excluded by KARMEN~\\cite{karmen} and awaits confirmation by MiniBooNE~\\cite{miniboone}. The current generation of nuclear reactor experiments (Palo Verde~\\cite{paloverde} and CHOOZ~\\cite{chooz}) find null oscillation results and rule out $\\bar{\\nu}_{e}\\rightarrow \\bar{\\nu}_x$ oscillations for $\\Delta m_{21}^2 \\gsim 10^{-3}$ $\\rm{eV}^2$ at maximal mixing and \\mbox{$\\rm{sin}^2 2\\,\\theta_{13} > 0.1$} for larger $\\Delta m_{21}^2$ (at the 95\\% confidence level).  The focus of the present work is the solar neutrino puzzle \\cite{bahcall} and KamLAND's~\\cite{kamland} role in resolving whether or not the large mixing angle (LMA) solution (see {\\it e.g.} Ref.~\\cite{lma}) is the correct one. The LMA solution has $\\Delta m_{21}^2$ in the range $10^{-5}$ to $10^{-4}$ $\\rm{eV}^2$ and mixing amplitude  $\\rm{sin}^2 2\\,\\theta_{12} > 0.3$. Data from SuperKamiokande \\cite{osaka} now favors the LMA solution over the small mixing angle, low and vacuum oscillation solutions (see Ref.~\\cite{kras} for a discussion of the various oscillation solutions). Solar neutrino measurements at SNO \\cite{sno} should discriminate between different oscillation solutions, but perhaps not in its first year of operation \\cite{bahcall2}. If the LMA solution is the correct one, KamLAND will provide a precise determination of the oscillation parameters.  KamLAND is unique in its potential as the first terrestrial experiment to explore the solar neutrino anomaly. It will provide a definitive test of the LMA solution by either ruling it out or by pinning down the values of $\\Delta m_{21}^2$ and $\\rm{sin}^2 2\\,\\theta_{12}$. KamLAND should be able to provide this answer in a year from the start of running in spring 2001 and then provide an accurate determination of the solar neutrino oscillation parameters by the end of the three years over which it is expected to take data for. Particularly significant will be its ability to precisely determine $\\Delta m_{21}^2$ in the LMA region because unlike the case of solar neutrinos, the $\\Delta m_{21}^2$-dependent contribution to oscillations of the reactor neutrinos will not suffer from averaging of the oscillation $L/E$ dependence.  If the LMA solution is correct, $\\Delta m_{21}^2$-dependent $CP$-violating effects \\cite{CP0} can be large enough to be tested at very long baseline accelerator-based experiments~\\cite{CP}. With KamLAND promising a precisely known value of  $\\Delta m_{21}^2$, it will provide essential information for future experiments studying  $CP$-violation in the neutrino sector.  In Section II we provide a brief overview of the KamLAND experiment and the oscillation hypothesis in which we will work. Section III will be devoted to details of our simulation of the experiment and the subsequent data analysis. We conclude in Section IV. ",
        "conclusions": "The results of our work are succinctly summarized in Figs.~\\ref{chi0}--\\ref{chi2}. Within one year of operation, KamLAND measurements of the reactor $\\bar{\\nu}_e$ flux will establish whether or not the LMA solution (\\mbox{$10^{-5} \\lsim \\Delta m_{21}^2 \\lsim 10^{-4}\\, \\rm{eV}^2$} and $\\rm{sin}^2 2\\,\\theta_{12}>0.3$) is correct. If it is, in a span of three years the accuracy with which $\\rm{sin}^2 2\\,\\theta_{12}$, and especially $\\Delta m_{21}^2$, will be determined is striking. KamLAND will give us $\\rm{sin}^2 2\\,\\theta_{12}$ to an accuracy of $\\pm 0.1$ and $\\Delta m_{21}^2$ to within a factor of 2 for $\\rm{sin}^2 2\\,\\theta_{12}>0.2$; $\\Delta m_{21}^2$ will be known to an accuracy of $\\pm 10$\\% for  $\\rm{sin}^2 2\\,\\theta_{12}>0.7$  (at the $2\\sigma$ level). However, no knowledge will be gleaned about $\\theta_{13}$. Our conclusions assume that the experiment employs a 1000 ton fiducial volume. If the fiducial volume is instead 600 tons, Fig.~\\ref{chi2} applies for data collected over 20 months instead of 12 and the contours of Fig.~\\ref{chi0} will be roughly twice as large at the end of three years. Details of systematic uncertainties and the accuracy with which the background shape is known will cause minor changes in our results, but the overall conclusions will remain intact.  If the LMA solution is found to be correct, the precision with which  $\\Delta m_{21}^2$ will be pinned down at KamLAND will prove to be very beneficial to studies of $CP$-violation in the lepton sector in long baseline experiments~\\cite{morecp,lfv}.   \\vspace{0.25in}"
    },
    "0011/astro-ph0011287_arXiv.txt": {
        "abstract": "We report the results of a systematic comparison between the vertically averaged model and the vertically explicit model of steady state, Keplerian, optically thick $\\alpha$-discs. The simulations have concerned discs currently found in three different systems: dwarf novae, young stellar objects and active galactic nuclei. In each case, we have explored four decades of accretion rates and almost the whole disc area (except the narrow region where the vertically averaged model has degenerate solutions). We find that the one layer approach gives a remarkably good estimate of the main physical quantities in the disc, and specially the temperature at the equatorial plane which is accurate to within 30 $\\%$ for cases considered. The major deviations (by a factor $\\lesssim 4$) are observed on the disc half-thickness. The sensitivity of the results to the $\\alpha$-parameter value has been tested for $0.001 \\le \\alpha \\le 0.1$ and appears to be weak. This study suggests that the ``precision'' of the vertically averaged model which is easy to implement should be sufficient in practice for many astrophysical applications. ",
        "introduction": "Recent investigations on MHD turbulence have confirmed the existence of a mechanism able to extract angular momentum from weakly magnetized accretion discs (Balbus \\& Hawley, 1991; Fleming, Stone \\& Hawley, 2000; Miller \\& Stone, 2000). Many aspects of accretion are, however, still not understood. In the meanwhile, present disc models mostly incorporate the $\\alpha$-prescription (Shakura \\& Sunyaev, 1973) and are of increasing complexity. For instance, models of steady state, Keplerian, vertically stratified $\\alpha$-discs (e.g. Cannizzo, 1992; Hameury et al., 1998; D'Alessio et al., 1999) should definitely supplant the more widespread, vertically averaged version (Pringle, 1981) which, despite a great simplicity and flexibility, cannot account for fine physical effects, remains in essence somewhat limited and contains some artefacts, like degenerate solutions (Collin-Souffrin \\& Dumont, 1990). Vertically explicit disc models offer in principle a much better reliability but, due to the uncertainty on the $\\alpha$-formalism, they require an extra prescription specifying the depth-dependence of the viscous energy release (Meyer \\& Meyer-Hofmeister, 1982; Mineshige \\& Osaki, 1983; Cannizzo \\& Cameron, 1988; Hameury et al., 1998). From a numerical point of view, the bi-dimensional problem which involves partial/ordinary differential equations is much less trivial to solve than the vertically averaged problem, also meaning a much larger computational time (by a factor $\\sim 100-1000$ say, corresponding to the typical resolution in the $z$-direction). It is therefore legitimate to ask, from a {\\it quantitative point of view} at least, which benefits the vertically explicit model really brings. The aim of this report is to show how, for steady Keplerian $\\alpha$-discs, the one layer approach and the vertically explicit approach compare. The main motivation is both to seek where the differences appear in an $\\alpha$-disc and to quantify these as functions of the accretion rate and value of the viscosity parameter. This investigation concerns three generic systems containing a disc: an active galactic nucleus, a young stellar object and a dwarf nova. \\begin{figure*}[htbp] \\psfig{figure=hure_galliano.f1,width=17.cm} \\caption{Midplane temperature ({\\it top}) and disc half-thickness ({\\it bottom}) versus the radius for $M=0.61$ M$_\\odot$, $\\dot{M}=10^{-10}$ M$_\\odot$/yr and $\\alpha=10^{-2}$ (DN; {\\it left}), $M=1$ M$_\\odot$, $\\dot{M}=10^{-7}$ M$_\\odot$/yr and $\\alpha=10^{-2}$ (YSO; {\\it middle}) and $M=10^8$ M$_\\odot$, $\\dot{M}=10^{-2}$ M$_\\odot$/yr and $\\alpha=0.1$ (AGN; {\\it right}). The systematic comparisons between the VA-model ({\\it dashed lines}) and the VE-model ({\\it solid lines}) concern the hot and cold branches, outside the zone where the VA-model has a multivalued solution (MVS). } \\label{courbe} \\end{figure*} ",
        "conclusions": "Given the ranges of parameter values (masses, accretion rates and viscosity parameter values) considered in this study, we can conclude that the vertically averaged model of steady state Keplerian $\\alpha$-discs is remarkably reliable, since it provides solutions very close to those obtained from the vertically explicit version. The midplane temperature appears as the quantity least sensitive to the vertical averaging (with a $30\\%$ deviation only) whereas the disc thickness seems the most affected (a factor $\\sim 4$ in some cases). Problems dealing with the disc irradiation (even with a global energy budget) should therefore account for vertical stratification since the disc response is sensitive to the local flaring angle and radial gradient (Kenyon \\& Hartmann, 1987).  This investigation demonstrates that the vertically averaged approach should be sufficient for many applications (e.g. Pringle, 1997; Burderi, King \\& Szuszkiewicz, 1998; Collin \\& Hur\\'e, 1999), in particular if the level of accuracy required on the disc structure is low or moderate. Note however that the tables given here can be used to ``correct'' the vertically averaged solution if a better precision is necessary. More basically, these tables show the typical error on the knowledge of the disc structure, that is, some methodology effect. It is interesting, but comforting, to see that the model is much more sensitive to the physics it contains (like turbulent viscosity, convection, opacities, etc.) than to the method used to derive the disc structure."
    },
    "0011/astro-ph0011078_arXiv.txt": {
        "abstract": "Optical and X-ray studies have established the prevalence of significant substructure in clusters of galaxies, indicating that clusters are young systems, and that recent major mergers have occurred in many clusters.  Numerical simulations show that sub-cluster merging will result in significant tidal forces on disk galaxies in accreting sub-groups, which is likely to lead to the transformation of spirals to S0s in clusters.  Simulations predict simultaneous moderate starbursts in gas-rich disk galaxies in clusters which show on-going or recent merger activity, as well as enhanced galaxy--galaxy interactions and mergers. Observational data from studies of spectral indicators, both of starburst activity in late-type galaxies in nearby clusters, and of poststarburst activity in early-type galaxies in nearby and intermediate redshift clusters, support this scenario. ",
        "introduction": "Both optical studies of sky-projected cluster galaxy distributions (e.g. Baier 1979; Geller \\& Beers 1982) and studies of cluster X-ray morphologies (e.g. Mohr, Fabricant \\& Geller 1993; Buote \\& Tsai 1996) reveal the presence of significant substructure in clusters of galaxies. Subgroup velocity statistics (e.g. Dressler \\& Shectman 1988; Bird 1994; Girardi et al. 1997; Solanes, Salvador-Sol\\'{e} \\& Gonz\\'{a}lez-Casado 1999) confirm these results, and are in substantial agreement in showing that some 30--40\\% of clusters have statistically significant substructure.  Numerical simulations also show that subclustering is a typical property of cluster formation in hierarchical theories of structure formation.  Such subclustering is expected for a broad range of cosmologies, and its predicted extent is in agreement with that observed (e.g. Knebe \\& M\\\"{u}ller 2000). Its prevalance indicates that clusters are young objects, since such substructure is expected to be smoothed out over a few cluster crossing times.  Moreover the existing substructure indicates that recent major mergers have occurred in many clusters.  The evolution of cluster substructure is potentially of great importance for its influence on the star formation histories and morphological evolution of cluster galaxies (e.g Tully \\& Shaya 1984; Lavery \\& Henry 1988; Gnedin 1999; Bekki 1999).  However, until very recently this potential influence has largely been neglected (cf. Bekki 1999). Of particular interest, is its relevance to the evolution of disk galaxies in clusters.  As is well known, Butcher \\& Oemler (1978) found that the cores of intermediate redshift clusters contain a higher fraction of blue, star-forming galaxies than similar environments at the present. These galaxies have been shown to be normal spiral or irregular galaxies, a fraction of which are obviously interacting or disturbed (e.g. Dressler et al. 1994; Oemler, Dressler \\& Butcher 1997; Smail et al. 1997).  In rich clusters they constitute up to 50\\% of the population, but by the present epoch have been depleted by a factor of two and replaced by a corresponding increase in S0s (e.g. Dressler et al. 1997).  Wang \\& Ulmer (1997) have found a correlation between the fraction of blue galaxies in intermediate redshift clusters and the cluster X-ray ellipticities.  They interpret the latter as an age sequence, and conclude that the younger clusters have a higher fraction of blue, star-forming galaxies. As the authors note, this interpretation is in accord with a hierarchical clustering scenario, in which a typical distant cluster is assembled from smaller units which tend to contain more gas-rich, late type galaxies.  The relative depletion of such star-forming galaxies in older, more relaxed clusters, suggests that the process of sub-cluster merging plays a crucial role in the transformation of an initial population of cluster spirals to a subsequent S0 population.  What effect is the sub-cluster merging process likely to have on a population of gas-rich, star-forming disk galaxies? Bekki (1999) using numerical simulation of a merger between a small group of galaxies and a cluster, shows that the time-dependent tidal gravitational field of the merger gives strong nonaxisymmetric perturbations to the disk galaxies in the group, inducing simultaneous moderate secondary starbursts in the central regions of the galaxies.  Similarly, Gnedin (1999) using self-consistent cluster simulations, demonstrates that for substantial infall of additional material into a cluster, the time-varying gravitational potential causes a sequence of tidal shocks on individual galaxies over a wide area of the cluster. These shocks enhance galaxy--galaxy interactions as well as amplifying galaxy merger rates.  Again the effect is also likely to induce simultaneous starbursts in gas-rich infalling disk galaxies over a wide area of the cluster. Moreover, the tidal shocks are shown to typically thicken the disks of large spirals by a factor of two, making it unlikely that spiral structure and gaseous shocks in the inner regions of the galaxies will form. A significant amount of the dark matter halo is stripped, possibly removing the resevoir of fresh gas which maintains the star formation activity (cf. Larson, Tinsley \\& Caldwell 1980).  While gas in the outer regions of the disk can be stripped by the ram pressure of the ambient medium, interpenetrating collisions with neighbouring galaxies could remove all their gas.  The effect of all these tidal transformations is the conversion of spirals to S0s. ",
        "conclusions": "Strong theoretical and observational data suggest that sub-cluster merging plays a crucial role in the transformation of spirals to S0s in the process of cluster formation, due to the strong tidal forces on the disk galaxy members of an accreting sub-group. Theoretical studies predict simultaneous moderate starbursts in disk galaxies of the accreting sub-group due to tidal forces.  This prediction is supported by observational data for both nearby and intermediate redshift clusters.  Tidal shocks due to the time-varying potential associated with sub-cluster merging are expected to effect morphological transformation of star-forming disk galaxies to S0s.  Thus these can explain the transformation of the majority of the spiral population in younger, less relaxed clusters to a population of S0s in older, more relaxed clusters as an integral part of the process of cluster formation."
    },
    "0011/astro-ph0011552_arXiv.txt": {
        "abstract": "Three observational constraints can be placed on a warm-hot intergalactic medium (WHIM) using \\rosat\\ PSPC pointed and survey data, the emission strength, the energy spectrum, and the fluctuation spectrum. The upper limit to the emission strength of the WHIM is $7.5\\pm1.0$ keV s$^{-1}$ cm$^{-2}$ sr$^{-1}$ keV$^{-1}$ in the \\tqkev\\ band, an unknown portion of which value may be due to our own Galactic halo. The spectral shape of the WHIM emission can be described as thermal emission with $\\log T=6.42$, although the true spectrum is more likely to come from a range of temperatures. The values of emission strength and spectral shape are in reasonable agreement with hydrodynamical cosmological models. The autocorrelation function in the 0.44 keV $< E <$1.21 keV band range, $w(\\theta)$, for the extragalactic soft X-ray background (SXRB) which includes both the WHIM and contributions due to point sources, is $\\lesssim0.002$ for $10\\arcmin < \\theta < 20\\arcmin$ in the \\tqkev\\ band. This value is lower than the \\citet{cmdhkfw2000} cosmological model by a factor of $\\sim5$, but is still not inconsistent with cosmological models. It is also found that the normalization of the extragalactic power law component of the soft X-ray background spectrum must be $9.5\\pm0.9$ keV s$^{-1}$ cm$^{-2}$ sr$^{-1}$ keV$^{-1}$ to be consistent with the \\rosat\\ All-Sky Survey. ",
        "introduction": "The problem of the missing baryons is well known; measurements of D in high redshift absorption line systems and application of Big Bang Nucleosynthesis models reveal a baryon fraction, $\\Omega h_{70}^2=0.0394\\pm0.0029$, \\citep{bt98} while inventories of local material can account for only $\\Omega h_{70}^{-1.5}=0.014\\pm0.0034$ \\citep{fhp98}. Cosmological hydrodynamical simulations (\\citealt{ckor95}, \\citealt{co99}, \\citealt{bn98}, \\citealt{dave2000}, and \\citealt{cmdhkfw2000} for example) suggest that the missing baryons exist as shock-heated low-density filaments of a ``warm-hot'' intergalactic medium (WHIM) at temperatures of $\\log T =5-7$. The visibility of the WHIM is compromised by a number of Galactic foregrounds that exist in the lower part of this temperature range, the Local Hot Bubble ($\\log T=6.11$, \\citealt{sealhb98}) and the Galactic Halo ($\\log T=6.08$ and, perhaps $\\log T=6.45$, \\citealt{ks2000}). All of these components emit predominately below E=1 keV, where Galactic absorption can be significant. A component at the high end of the WHIM range (e.g., $\\log T=7$) would emit strongly in the 1.5 keV band. However, \\citet{hea98} has directly resolved 70-80\\% of the flux in the 0.5-2.0 keV range, leaving little room for diffuse components. Thus, significant diffuse emission at the high end of the WHIM range would be difficult to reconcile with observations, while emission at the low end of the WHIM range becomes difficult to detect.  Without high resolution, high grasp, non-dispersive X-ray spectrometers, there are three observational tests of the WHIM models; the total emission, the energy spectrum, and the fluctuation spectrum. Theoretical calculations by \\citet{co99} suggest that the WHIM produces $\\sim7$ keV s$^{-1}$ cm$^{-2}$ sr$^{-1}$ keV$^{-1}$ at 0.7 keV; the \\citet{cmdhkfw2000} prediction for the entire IGM in the 0.5-2.0 keV band yields the same value for the \\tqkev\\ band (\\rosat\\ R45 band). \\citet{poc2000}, however, estimate a substantially lower X-ray emission, $\\sim .7$ keV s$^{-1}$ cm$^{-2}$ sr$^{-1}$ keV$^{-1}$ in the 0.5-2.0 keV band. \\citet{dave2000} using a number of models including that of \\citet{cmdhkfw2000} predicts that the baryon fraction as a function of temperature will be peaked near $\\log T=6.45$. Finally, \\citet{cmdhkfw2000} predict that the value of the ACF of the entire extragalactic SXRB in the 0.5-2.0 keV band should be $\\sim0.005$ for $\\theta=10\\arcmin$. ",
        "conclusions": "The observed emission strength, energy spectrum, and fluctuation spectrum of the soft X-ray background originating from beyond the Galactic absorption is roughly consistent with the predictions by hydrodynamical simulations of the universe. Even given the caveats concerning the Galactic halo, the current observations do not {\\it disprove} the existing models. The observations here present solid guideposts for further exploration of the accessible parameter space by future generations of simulations.  We wish to thank J.P. Ostriker and the anonymous referee for their very useful comments."
    },
    "0011/astro-ph0011308_arXiv.txt": {
        "abstract": "\\PSR\\ is a 47ms radio pulsar with a high spin-down luminosity which is in a close, highly eccentric 3.5-year orbit about a bright stellar companion. The binary system may be a detectable source of hard $\\gamma$-rays produced by inverse Compton scattering of photons from the B2e star SS2883 by electrons and positrons in the pulsar wind. The star provides an enormous density of optical photons in the vicinity of the pulsar, particularly at epochs near periastron. We calculate the emission from the unshocked region of the pulsar wind, assuming that it terminates at a shock where it attains pressure balance with the companion's wind. The spectra and light curves for the inverse Compton emission from the shock-terminated wind are compared with those for an unterminated wind. If the pulsar's wind is weaker than that from the companion star, the termination of the wind decreases the inverse Compton flux, particularly near periastron. The termination shock geometry has the effect of decreasing the asymmetry of the $\\gamma$-ray light curve around periastron, which arises because of the asymmetrical variation of the scattering angle. ",
        "introduction": "\\PSR\\ is one of only three known radio pulsars which have a main sequence star binary companion. Such systems provide a unique environment for inverse Compton scattering because of the presence of an enormous density of low energy (optical) photons which serve as targets for electrons and positrons in the pulsar wind. The modest spin-down luminosity, and relatively high distance to two of these systems together imply that they are unlikely to be detectable sources of inverse Compton emission. In contrast, \\PSR\\ is a galactic radio pulsar with a high spin-down luminosity ($\\Lp = 8.3\\times 10^{28}\\;$W) and is relatively nearby ($1.5\\,$kpc). If just 0.1\\% of the wind luminosity is scattered into hard $\\gamma$-ray photons the resulting flux should be detectable using current $\\gamma$-ray telescopes.  \\PSR\\ is in a highly eccentric orbit ($e\\sim 0.87$) around SS2883, a B2e star of radius $\\Rstar\\sim 6 R_\\odot$ and luminosity $\\Lstar\\sim 8.8\\times 10^3\\,L_\\odot$ [Johnston et al.\\ 1992, 1994, 1996]. At periastron the pulsar is only $23\\Rstar\\approx 10^{11}\\,{\\rm m}$ from its companion. Around periastron there are $3.4\\times 10^{11}\\,\\rm cm^{-3}$ Be-star photons at the pulsar, corresponding to an energy density of $\\Urad\\sim 6\\times 10^{11}\\,\\rm eV\\,cm^{-3}$. This is some 11 orders of magnitude larger than the typical background target density available for inverse Compton scattering by the winds of isolated pulsars.  The Crab is presently the only pulsar for which observations place any constraints on the physical parameters of a pulsar wind. Observations of the \\PSR\\ system aimed at detecting hard $\\gamma$-ray inverse Compton emission are planned during 2000. If successful they should provide the first direct probe of the freely-expanding region of the wind of any rotation-powered pulsar, and serve to constrain wind models and parameters.  In the absence of an efficient deceleration process the relativistic wind of a pulsar will expand freely until it attains pressure balance with the surrounding medium. If the wind is still supersonic at the point of pressure balance it will be bounded by a termination shock. Kirk, Ball \\& Skj\\ae raasen [1999] investigated the $\\gamma$-ray emission from \\PSR\\ resulting from inverse Compton scattering by the shocked pulsar wind, downstream of the termination shock, where the pulsar wind electrons and positrons have been accelerated and isotropised as the radial flow of the wind is disrupted. The effects of the cooling of the shocked wind by inverse Compton scattering were included by Tavani \\& Arons [1997] in their model for the synchrotron emission from the system. Ball \\& Kirk [2000a] considered the effects of inverse Compton scattering by the unshocked pulsar wind, upstream of the termination shock. They calculated the deceleration of the wind that occurs as a result of `inverse Compton drag' as energy and momentum are transferred from the wind to the scattered photons. In particular it was shown that the inverse Compton losses were unlikely to contain the wind of \\PSR\\ before it attains the radius at which pressure balance with the Be-star outflow is likely to occur. Ball \\& Kirk [2000a] then calculated the emission from the freely-expanding wind subject to the assumption that the wind was not terminated by pressure balance.  In this paper we recalculate the inverse Compton emission from the unshocked wind of \\linebreak\\PSR\\ including the effects of termination of the radially-expanding pulsar wind as a result of pressure balance with the Be-star wind. The relative strengths of the pulsar and Be-star winds are unknown. If the Be-star wind dominates then the termination shock will be close to, and wrapped around, the pulsar. This will have the effect of decreasing the emission from the unshocked portion of the wind, and the effects vary over the binary orbit of the system. We hope that inverse Compton $\\gamma$-rays will be detected from \\PSR\\ in the near future, providing data that can be compared with such models.  The calculation of the geometry of the termination shock is outlined in \\S 2. In \\S 3 we present $\\gamma$-ray spectra and light curves from the terminated wind, and compare them with those calculated for the unterminated wind by Ball \\& Kirk [2000a]. Our conclusions are presented in \\S 4. ",
        "conclusions": "Ball \\& Kirk [2000a] argued that inverse Compton emission from the freely-expanding portion of the wind of \\PSR\\ should be detectable at energies somewhere between GeV and TeV, depending on the wind parameters. The results presented in Section 3 suggest that even if the wind of the Be star dominates that of the pulsar, terminating it in a shock which wraps around the pulsar, the scattered emission should still be detectable for a wide range of wind parameters. However, the TS has the effect of decreasing the peak integrated $\\gamma$-ray flux from the freely-expanding wind, and of decreasing the asymmetry which otherwise produces a characteristic light curve which shows a higher $\\gamma$-ray flux before periastron than after.  Even if the Be-star wind dominates to such a degree that the unshocked pulsar wind is not detectable, the models of Kirk, Ball \\& Skj\\ae raasen [1999] indicate that inverse Compton emission from the shocked pulsar wind should be above the thresholds of current detectors.  The scattered $\\gamma$-ray signal from this system is most likely to comprise a combination of emission from the shocked and unshocked regions of the pulsar wind. The reduced asymmetry and periastron to apastron ratio of the emission from the terminated wind may make it harder to deconvolve the two contributions, but only detailed modelling of real data will provide such answers.  The best opportunity for detecting such emission in the near future is afforded by the new CANGAROO II imaging Cherenkov telescope located in Australia [Yoshikoshi et al.\\ 1999]. When the present lack of operational $\\gamma$-ray telescopes is relieved by the launch of the GLAST and INTEGRAL observatories, inverse Compton emission from this pulsar system may well be detectable at lower energies in the MeV--GeV range.  The next periastron of \\PSR\\ occurs in October 2000. CANGAROO II observations of the system will not be possible at times near periastron because at that time the object will be too close to the Sun. In any event, the light curves shown in Figures \\ref{light.1} and \\ref{light.01} are not realistic between about days $-25$ and $+25$ because we have not modelled the interaction with the Be-star disk which is important at such epochs. Observations of the \\PSR\\ system with CANGAROO II are planned for July and December 2000, when our models of the inverse Compton emission from the system should be applicable. If such observations are successful in detecting hard $\\gamma$-ray emission from this unique pulsar system they should provide valuable insights into the properties of the pulsar wind. In particular, detection of inverse Compton $\\gamma$-ray emission from the system should help to constrain values of the bulk Lorentz factor of the pulsar wind and the fraction of the spin-down luminosity carried by electrons and positrons compared to ions."
    },
    "0011/astro-ph0011414_arXiv.txt": {
        "abstract": "Here an attempt is made at summarizing the presentations, most of which were about the highest energy particles observed in nature. Particular attention is paid to the solutions to the Ultra High Energy Cosmic Ray particles, to the new and forthcoming data, to the new proposals for experiments and to the role of primary composition, that were amongst the most discussed subtopics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011179_arXiv.txt": {
        "abstract": "We report the measurement of the primordial D/H abundance ratio towards QSO \\object.  The column density of the neutral hydrogen in the $z \\simeq 2.536$ Lyman limit system is high, \\lnhi\\ $= 19.422 \\pm 0.009$ \\cmm , allowing for the deuterium to be seen in 5 Lyman series transitions. The measured value of the D/H ratio towards QSO \\object\\ is found to be D/H$ = 2.54 \\pm 0.23 \\times 10^{-5}$. The metallicity of the system showing D/H is found to be $\\simeq 0.01$ solar, indicating that the measured D/H is the primordial D/H within the measurement errors. The gas which shows D/H is neutral, unlike previous D/H systems which were more highly ionized.  Thus, the determination of the D/H ratio becomes more secure since we are measuring it in different astrophysical environments, but the error is larger because we now see more dispersion between measurements. Combined with prior measurements of D/H, the best D/H ratio is now D/H$ = 3.0 \\pm 0.4 \\times 10^{-5}$, which is 10\\% lower than the previous value. The new values for the baryon-to-photon ratio, and baryonic matter density derived from D/H are $\\eta = 5.6 \\pm 0.5 \\times 10^{-10} $ and \\ob $=0.0205 \\pm 0.0018$ respectively. ",
        "introduction": "The standard theory of big bang nucleosynthesis (SBBN) predicts the abundances of the light nuclei H, D, \\het , \\hef , and \\lisv\\ as a function of the cosmological baryon to photon ratio, $\\eta = {n_{b}}/{n_{\\gamma}}$ (Kolb \\& Turner 1990; Walker et al. 1991; Schramm \\& Turner 1998; Nollett \\& Burles 2000; Olive, Steigman \\& Walker 2000). A measurement of the ratio of any two primordial abundances gives $\\eta$, and hence the baryon density, while a second ratio tests the theory.  However, it is extremely difficult to measure primordial abundances, because in most places gas ejected from stars has changed the abundances.  Adams (1976) suggested that it might be possible to measure the primordial D/H ratio in absorption line systems towards QSOs. Although gas which has been inside a star will have lost all of its deuterium, in QSO absorption line systems having typical metal abundances of 0.001 to 0.01 of the solar value, about 0.1 -- 1\\% of the deuterium will have been lost.  The advent of the HIRES spectrograph (Vogt 1994) on the W.M. Keck-I telescope gave the high signal-to-noise ratio (SNR) and spectral resolution needed to reveal deuterium (Tytler et al. 2000) in high redshift absorption systems.  We have previously measured D/H in two QSOs (Tytler, Fan \\& Burles 1996; Tytler \\&\\ Burles 1997; Burles \\&\\ Tytler 1998a; Burles \\&\\ Tytler 1998b), and placed a strong upper limit on D/H in a third (Kirkman et al. 1999).  Other QSOs give less useful constraints on D/H, because their absorption systems are more complex, or existing spectra are inadequate. The Lyman limit system (LLS) at $z_{abs}=0.701$ towards QSO \\qhst\\ might allow ten times larger D/H, or it may give no useful constraints (Webb et al. 1997; Levshakov, Kegel, \\& Takahara 1998; Tytler et al. 1999). Molaro et al. (1999) claimed another QSO absorption system showed low D/H, but they and Levshakov et al. (2000) note that since only the \\lya\\ line has been observed, the hydrogen velocity structure and the H~I column density are poorly known, and the deuterium feature can be fit using hydrogen alone.  In this paper, we present a fourth QSO, \\object , which gives strong constraints on the primordial D/H ratio. ",
        "conclusions": "We have presented the fourth quasar, \\object , which shows an absorption system having a low deuterium to hydrogen abundance ratio: D/H$ = 2.54 \\pm 0.23 \\times 10^{-5}$ in the $z \\simeq 2.536$ Lyman limit absorption system.  We first obtained low resolution spectra in our survey for D/H QSOs with the Lick 3 meter Shane telescope, and here we presented over 24 hours of spectra from the HIRES spectrograph on the Keck-I telescope.  The absorber has a high neutral hydrogen column density, \\lnhi $=19.422 \\pm 0.009$ \\cmm\\ which is 36 times larger than the next highest case studied for D/H, but 8--100 times less than standard damped Lyman alpha systems. Very little is known about absorbers with \\lnhi $\\simeq 19$ \\cmm , and this absorber may not be representative of this class because it was selected to have a very simple velocity structure and low $b$ values.  While the absorber towards object has by far the highest neutral H~I column density of the absorbers which we have studied for D/H, it also has the lowest total Hydrogen column density, by a factor of 4 -- 7, when we correct for the ionization.  We know little about the environment around the absorber. It might be in the outer parts of galaxy, as are LLS at low redshift, or in a disk, as are damped Lyman alpha absorbers at high redshift. Alternatively, it may be in a relatively isolated gas cloud, because see just one component, with an exceptionally small spread of velocities. In either case, numerical simulations of the growth of structure suggest that the absorbing gas has been incorporated into a galaxy by today.  For the first time, we detect deuterium absorption in 5 Lyman series transitions, and determine \\lndi $=14.826 \\pm 0.039$ \\cmm .  We have strong arguments that the observed absorption is indeed deuterium, and not inter-loping hydrogen or metal line absorption.  We observe a number of associated metal line absorbers, from which we calculate that the gas is warm and neutral. The metallicity of the system is $\\simeq 0.01$ times solar, indicating that the measured D/H is representative of primordial D/H.  We argue that \\object\\ offers the most secure detection of D/H to date, and that the D/H ratio determined from all QSOs has been made more secure.  \\subsection{\\object\\ gives the most secure measurement of primordial D/H}  The measurement of primordial D/H towards \\object\\ is more secure than our prior measurements towards \\qfirst , \\qsecond\\ and \\qthird\\ for several reasons. By secure we mean that we have the most information, and hence there is less chance of undetected errors which might greatly exceed those quoted. We are not explicitly referring to the size of the quoted errors, which are similar for \\qfirst\\ and \\object , and larger for \\qsecond\\ .  The absorption system in \\object\\ is simple. Like \\qthird , the second most secure result, the absorber in \\object\\ is modeled with a single component, which simplifies the measurement of the column densities. \\qfirst\\ was modeled with two or three components, and \\qsecond\\ with two to four.  \\object\\ is the only case to show more than one strong deuterium line, which reduces the chance of contamination, and gives more reliable $b$ values, and hence \\ndi .  For \\object\\ we listed above the many reasons why the absorption near the deuterium line position is deuterium. The chance of serious contamination from the \\lya\\ forest will decrease with rising \\ndi . Such contamination is least likely to be significant in \\object , followed by \\qfirst . There is some contamination in \\qsecond\\ while for \\qthird\\ we see a lot of contamination, and obtain only an upper limit on \\ndi . We do not know whether the chance of contamination by components of the Lyman limit system changes with \\nhi .  We also expect that the chance of contamination decreases as the $b$ value of the deuterium decreases, because H lines often have larger $b$ values. Hence \\object\\ is the most secure detection of deuterium.  For \\object\\ we have the most information on the velocity field and $b$ values because we see several D~I, N~I and O~I lines.  For \\object\\ the metal abundance is obtained with additional redundancy. The gas is nearly neutral, and hence we get the abundances of several elements: H, C, N, O, Al, Fe and Si. For the \\qfirst\\ and \\qsecond\\ we used a standard photoionization model to find the level of ionization which explained the relative abundances of ions such as C~II, C~III, C~IV and Si~II, Si~III, Si~IV. We obtained a solution for each element, and these agreed, which provided a check. For \\qthird\\ the ionization and metal abundances are both less well known.  \\subsection{The primordial D/H becomes more secure}  The new measurement makes the primordial D/H much more secure because in each case we are sampling gas with different physical conditions, and some systematic errors, including those associated with the measurement of column densities, may be different for each QSO. First, the new measurement is the most secure. Second, we have increased the number of QSOs in which we have measured deuterium from two to three. Third, we sample a new region of space. Each absorber is a different direction in the universe, and each samples a sight line of about 1 -- 10~kpc, which requires about 10$^5 - 10^6$ solar masses of gas. Fourth, the absorbing gas covers a factor of 240 range in \\nhi . Fifth, the absorption systems have differing ionization, with a range of 2000 in the H~I/H ratio.  There is less variation in other parameters. The metal abundances cover a factor of ten, which is the typical for QSO absorption line systems, while the redshifts cover most of the range observable from the ground."
    },
    "0011/astro-ph0011453_arXiv.txt": {
        "abstract": "We describe automatic procedures for the selection of DA white dwarfs in the Hamburg/ESO objective-prism survey (HES). For this purpose, and the selection of other stellar objects (e.g., metal-poor stars and carbon stars), a flexible, robust algorithm for detection of stellar absorption and emission lines in the digital spectra of the HES was developed. Broad band ($U-B$, $B-V$) and narrow band (Str\\\"omgren $c_1$) colours can be derived directly from HES spectra, with precisions of $\\sigma_{U-B}=0.092$\\,mag; $\\sigma_{B-V}=0.095$\\,mag; $\\sigma_{c_1}=0.15$\\,mag.  We describe simulation techniques that allow to convert model or slit spectra to HES spectra. These simulated objective-prism spectra are used to determine quantitative selection criteria, and for the study of selection functions. We present an atlas of simulated HES spectra of DA and DB white dwarfs.  Our current selection algorithm is tuned to yield maximum efficiency of the candidate sample (minimum contamination with non-DAs). DA candidates are selected in the $B-V$ versus $U-B$ and $c_1$ versus $W_\\lambda(\\mbox{H}\\beta +\\mbox{H}\\gamma+\\mbox{H}\\delta)$ parameter spaces. The contamination of the resulting sample with hot subdwarfs is expected to be as low as $\\sim 8$\\,\\%, while there is essentially no contamination with main sequence or horizontal branch stars. We estimate that with the present set of criteria, $\\sim 80$\\,\\% of DAs present in the HES database are recovered. A yet higher degree of internal completeness could be reached at the expense of higher contamination. However, the external completeness is limited by additional losses caused by proper motion effects and the epoch differences between direct and spectral plates used in the HES. ",
        "introduction": "The Hamburg/ESO survey (HES) is an objective prism survey primarily targeting bright quasars \\citep{hespaperI,heshighlights,hespaperIII}. However, because its spectral resolution is typically 15\\,{\\AA} FWHM at H$\\gamma$, it is also possible to efficiently select a variety of interesting \\emph{stellar} objects in the HES. These include, e.g., metal-poor halo stars, carbon stars, cataclysmic variable stars, white dwarfs (WDs), subdwarf B stars (sdBs), subdwarf O stars (sdOs), and field horizontal branch A- and B-type stars \\citep{Christlieb:2000}.  In a series of papers, we will report on the development of quantitative selection procedures for the systematic exploitation of the stellar content of the HES, and their application to the digital HES data base. In this paper, we report on an automatic algorithm for the selection of DA white dwarfs (DAs).  \\begin{figure*}[htbp] \\begin{center} \\epsfig{file=ms10235f1.ps, clip=, width=14cm, bbllx=48, bblly=420, bburx=527, bbury=740} \\caption{Definition of HES area (framed) and numbers of fields in which the exploitation of the stellar content of the HES is currently carried out. For orientation, the position of the southern galactic pole is marked with `$\\ast$' (lower left corner of field 475). } \\end{center} \\end{figure*}  The aim of the DA selection in the HES described in this paper is to test the double-degenerate (DD) scenario for SN~Ia progenitors, in which a binary consisting of two WDs of large enough mass, merges and produces a thermonuclear explosion. If this scenario is correct, SN~Ia progenitor systems should be present among DDs, the latter being identifiable by radial velocity (RV) variations. Although several double degenerates have been found, none of them is sufficiently massive to qualify as a viable SN Ia progenitor \\citep[see][]{Maxted/Marsh:1999}. Note however that recently a system consisting of a sdB and a massive white dwarf has been found, the total mass of which exceeds the Chandrasekhar mass \\citep{Maxtedetal:2000a}. In order to increase the sample of DDs, a \\emph{Large Programme} was proposed to (and accepted by) the European Southern Observatory (ESO), aiming at observing a large ($\\sim 1\\,500$) sample of WDs with VLT Unit Telescope 2 (UT2), and its high resolution spectrograph UVES, at randomly chosen epochs, dictated by observing conditions. That is, every time the weather is \\emph{too bad} to carry out Service Mode observations for other programs, WDs are observed.  A program like this requires a large catalog of targets spread all over the accessible sky in order to be successfull. From the ``Catalog of Spectroscopically Identified White Dwarfs'' of \\cite{McCook/Sion:1999} it is evident that the southern sky has not yet been surveyed as extensively for white dwarfs as the northern sky: only 33\\,\\% of the objects listed are located at $\\delta<0^{\\circ}$. Therefore, we were aiming at selecting additional targets in the data base of the HES.  White dwarfs have been selected from wide angle surveys in the southern hemisphere before, and also in the HES (see below). ``Classical'' UV excess surveys, like the Montreal-Cambridge-Tololo survey \\citep[MCT; ][]{Demersetal:1986,Lamontagneetal:2000}, or the Edinburgh-Cape survey \\citep[EC; ][]{Stobieetal:1997,Kilkennyetal:1997} can efficiently select complete samples of hot stars, including WDs. However, completeness at the \\emph{cool} end is either sacrificed for efficiency, as in the MCT (see Fig. \\ref{UBV_DA}), where only objects with $U-B<-0.6$ enter the sample of stars for which follow-up spectroscopy is obtained \\citep{Lamontagneetal:2000}, or efficiency is sacrificed for completeness, as in the EC. It has been shown that the EC is 94\\,\\% complete for objects of $U-B<-0.4$ down to $B=16.5$ \\citep{Stobieetal:1997}. However, an intermediate selection step based on photoelectric $UBV$ photometry has to be used to eliminate the large fraction ($\\sim 30$\\,\\%; see \\citealt{Kilkennyetal:1997}) of ``normal'' F and G type stars.  \\begin{figure*}[htbp] \\begin{center} \\leavevmode \\epsfig{file=ms10235f2.ps, clip=, width=14cm, bbllx=67, bblly=250, bburx=460, bbury=514} \\caption{ \\label{UBV_DA} $U-B$ and $B-V$ for DAs as a function of effective temperature, determined with model spectra. Taking into account the r.m.s. error in $U-B$ of 0.18\\,mag in the MCT survey, the adopted selection criterion of $U-B<-0.6$ leads to rejection of cool ($T_{\\mbox{\\scriptsize eff}} \\lesssim 22\\,000$\\,K) DA white dwarfs. In the HES, only DAs of $T_{\\mbox{\\scriptsize eff}}\\lesssim 11\\,000$\\,K are expected to be lost when the selection criteria described in Sect. \\ref{Sect:DAsel} are applied.} \\end{center} \\end{figure*}  In the HES, WDs enter the quasar candidate sample if they have $U-B<-0.18$ \\citep{hespaperIII}. However, HES quasar candidates are inspected manually at the computer screen, and in this process hot stars, and stars clearly exhibiting stellar absorption lines (like e.g. DA white dwarfs, having strong, broad lines over a wide temperature range; see Fig. \\ref{WDmodels} in Appendix \\ref{Sect:SpectralAtlas}) are rejected, and follow-up spectroscopy is not obtained for them in the course of the quasar survey. This results in a very efficient quasar selection: typically 70\\,\\% of the objects for which follow-up spectroscopy is obtained \\emph{are} quasars \\citep{hespaperIII}. The remaining 30\\,\\% are mainly hot subdwarfs, cool ($T_{\\mbox{\\scriptsize eff}}\\lesssim 20\\,000$\\,K), helium-rich WDs \\citep[DBs, DZs; see][]{Friedrichetal:2000}, a couple of interesting peculiar objects, e.g., magnetic DBs \\citep{magDB1} and magnetic DAs \\citep{Reimersetal:1994,Reimersetal:1996} have also been discovered in this way.  The selection of white dwarf candidates described in this paper aims at an \\emph{efficient} selection; that is, the contamination of the sample with other objects (e.g., hot subdwarfs) shall be as clean as possible, in order not to waste any observing time at the VLT. ",
        "conclusions": "We described quantitative procedures for the selection of DA white dwarfs in the digital data base of the HES. Algorithms for the detection of stellar emission and absorption lines, broad- and intermediate-band colours were developed. These are not only used for the selection of DAs, but also for a couple of other interesting stars, e.g. metal-poor stars, and carbon stars. Simulation techniques allow us to convert model spectra to HES spectra, which is important for the development of selection criteria, for the compilation of learning samples used for automatic spectral classification, and for the determination of selection functions.  DAs can be selected very efficiently in the HES, as required by a currently ongoing \\emph{Large Programme} using VLT-UT2 and UVES, aiming at the detection of \\emph{double degenerates}. A first set of 440 DA candidates not listed in the catalog of \\cite{McCook/Sion:1999} was identified on 225 HES plates (i.e., 59\\,\\% of the survey).  We note that the UVES sample includes a double-lined DA$+$DA binary, and two further DAs with significant RV variations. The results will be reported elsewhere. Further investigation of these systems is in progress.  For investigations requiring \\emph{complete} samples of DAs (or other objects having high proper motions), special extraction techniques have to be employed to take into account the epoch differences between HES plates and the DSS~I, since the latter is used for object detection in the HES. Such techniques are currently under development."
    },
    "0011/astro-ph0011386_arXiv.txt": {
        "abstract": "A remarkably precise observational relation for pulse core component widths of radio pulsars is used to derive stringent limits on pulsar radii, strongly indicating that pulsars are strange stars rather than neutron stars. This is achieved by inclusion of general relativistic effects due to the pulsar mass on the size of the emission region needed to explain the observed pulse widths, which constrain the pulsar masses to be $\\le \\; 2.5 \\; M_\\odot$ and radii $\\le \\; 10.5 \\;km$. ",
        "introduction": "Radio pulsars are believed to be the most common manifestations of neutron stars, but it has not been possible so far to relate the voluminous data on radio pulses and their varied structure to the properties of neutron stars except through the arrival times of pulses. Here we make such a connection between pulse core component widths derived from very good quality radio data and the mass-radius (M-R) relation of neutron stars. This becomes possible only due to the inclusion of general relativistic effects of the stellar mass on pulsar beam shapes, which makes the stellar mass and radius relevant parameters in determining the pulse widths. We show that core component widths provide tight constraints on equations of state (EOS) of neutron stars. We compare our results with other similar attempts based, e.g., on the X-ray data. From our constraints it emerges that no neutron star EOS seem to be adequate, leading to the conclusion that pulsars are strange stars, i.e., ones composed of quarks of flavors u, d, and s (Alcock et al. 1986); and we examine it in light of similar recent suggestions. ",
        "conclusions": "In summary, the empirical formula of Rankin (1990) describing the opening angle of the pulsar beam emitting the core emission when compared to theoretically calculated value leads to a constraint that pulsar masses should be $\\le \\; 2.5 \\; M_\\odot$ and radii $\\le \\; 10.5 \\; km$. This comes about due to the inclusion of general relativistic effects of the mass of the star on the pulsar beam size. For observationally reasonable pulsar masses a comparison with mass-radius relations of neutron star EOS shows that most of the EOS are ruled out, implying that pulsars are strange stars and not neutron stars, unless our understanding of the neutron star EOS is revised.  {\\it Acknowledgements} One of us (C.S.S.) would like to thank Rajeev Gavai and Sanjay Reddy for rekindling his interest in the subject matter of this paper whose preliminary version was reported in Kapoor and Shukre (1997). We thank I. Bombaci and Joanna Rankin for valuable comments."
    },
    "0011/gr-qc0011078_arXiv.txt": {
        "abstract": "The radion on the de Sitter brane is investigated at the linear perturbation level, using the covariant curvature tensor formalism developed by Shiromizu, Maeda and Sasaki\\cite{SMS}. It is found that if there is only one de Sitter brane with positive tension, there is no radion and thus the ordinary Einstein gravity is recovered on the brane other than corrections due to the massive Kaluza-Klein modes. As a by-product of using the covariant curvature tensor formalism, it is immediately seen that the cosmological scalar, vector and tensor type perturbations all have the same Kaluza-Klein spectrum. On the other hand, if there are two branes with positive and negative tensions, the gravity on each brane takes corrections from the radion mode in addition to the Kaluza-Klein modes and the radion is found to have a negative mass-squared proportional to the curvature of the de Sitter brane, in contrast to the flat brane case in which the radion mass vanishes and degenerates with the $4$-dimensional graviton modes. To relate our result with the metric perturbation approach, we derive the second order action for the brane displacement. We find that the radion identified in our approach indeed corresponds to the relative displacement of the branes in the Randall-Sundrum gauge and describes the scalar curvature perturbations of the branes in the gaussian normal coordinates around the branes. Implications to the inflationary brane universe are briefly discussed. ",
        "introduction": "Based on the idea of a brane-world suggested from string theory\\cite{Antoniadis}, Randall and Sundrum proposed an interesting scenario that we may live on a 3-brane embedded in the 5-dimensional anti-de Sitter space\\cite{RS1,RS2}. One of the attractive features of their scenario is that gravity on the brane may be confined within a short distance from the brane to be effectively 4-dimensional even for an infinitely large extra-dimension\\cite{RS2,GT}. This occurs on a brane with positive tension, and it is because the AdS bulk on both sides of the brane shrinks exponentially as one goes away from the brane.  In their comprehensive paper\\cite{GT}, Garriga and Tanaka has shown that if there is only one flat brane with positive tension, the zero-mode truncation, which means to discard all the excitation modes except for the one with the lowest (zero) eigenvalue, leads to 4-dimensional (linearized) Einstein gravity on the brane. On the other hand, if there are two flat branes of positive and negative tensions, the zero-mode truncation leads to a Brans-Dicke type theory with positive and negative Brans-Dicke parameters, respectively, on the branes with positive and negative tensions. The difference can be understood by considering the role of the so-called radion mode, or the unique non-trivial scalar mode of gravitation in the vacuum bulk spacetime. (Hereafter, we use the term, scalar, in the sense of the 4-dimensions unless otherwise noted.) The radion mode corresponds to displacement of the brane location\\cite{GT,CGR}. If only a single flat brane is present, it is impossible to know locally if the brane is displaced because of the general covariance. Hence, the radion plays no role in the gravity on the brane. However, if two flat branes are present, the relative motion of the branes, or the relative displacement of the branes is physical. Therefore the radion mode is expected to affect the gravity on the brane. Thus if we live on a single positive-tension brane, the effect of the extra-dimension appears only through the non-zero-mode excitations, or the so-called Kaluza-Klein modes.  However, the above interpretation is given to the case of flat Minkowski branes. It is therefore interesting to ask if the same argument applies to the case of cosmological brane models, particularly when the universe is spatially closed. It was suggested that a spatially closed, inflating brane-universe may be created from nothing\\cite{GS}. In this model, the universe is described by the de Sitter brane embedded in the 5-dimensional anti-de Sitter space. Since this brane-universe is spatially closed, the curvature radius of the universe becomes indeed the 5-dimensional radius of the brane. Then, one may expect that the radion, which describes fluctuations of the brane location, becomes physical even in the case of a single brane, and contributes to gravity on the brane.  In this paper, we carefully investigate perturbations of the de Sitter brane-world and clarify this issue of the cosmological radion. Cosmological brane perturbations have been studied by various authors with metric-based approaches\\cite{GS,Mukohyama,Kodama,HHR,Carsten,Koyama,Langlois,% MWBH,LMW,bmw,cll,DDK,NeroSachs}. We adopt the covariant curvature tensor formalism developed by Shiromizu, Maeda and Sasaki\\cite{SMS}. A covariant approach to cosmological brane perturbations based on this formalism has been developed by Maartens\\cite{Roy}. In this formalism, the effective gravitational equation is written in the form of the 4-dimensional Einstein equations with a couple of modifications on the right-hand-side. Namely, there appear a tensor quadratic in the matter energy-momentum tensor and an ``electric'' projection of the 5-dimensional Weyl tensor. Since the former is locally described by the 4-dimensional energy-momentum tensor, the genuine 5-dimensional gravitational effect is contained in the projected Weyl tensor. We denote this tensor by $E_{\\mu\\nu}$ (see Eq.~(\\ref{Edefined}) for its definition).  A noble feature of this formalism is that the perturbations that describe the 4-dimensional gravitational degrees of freedom are already contained in the effective 4-dimensional Einstein equations and the perturbations that additionally arise from the existence of the 5-dimensional bulk are completely described by the projected Weyl tensor mentioned above. Another advantage of the formalism is that since the 5-dimensional Weyl tensor vanishes on the anti-de Sitter background, $E_{\\mu\\nu}$ is a physical, gauge-invariant quantity by itself.  In this paper we neglect the tensor quadratic in the matter energy-momentum and investigate the behavior of the projected Weyl tensor in detail at the linear perturbation level. We assume the AdS bulk with the de Sitter brane as the background. On this background, the 5-dimensional Weyl tensor vanishes, hence $E_{\\mu\\nu}=0$. We then assume the matter energy momentum tensor on the brane, $\\tau_{\\mu\\nu}$ to be small and consider $E_{\\mu\\nu}$ induced by $\\tau_{\\mu\\nu}$.  This paper is organized as follows. In Sec.~II, we derive the general solution to the projected 5-dimensional Weyl tensor perturbation, $E_{\\mu\\nu}$. In Sec.~III, we concentrate on the cosmological radion mode and derive the explicit form of $E_{\\mu\\nu}$ in terms of the radion after zero-mode truncation. The physical meaning of the radion is also discussed. In Sec.~IV, we summarize the result and discuss its implications. ",
        "conclusions": "In this paper, we have investigated the role of the radion mode on the de Sitter brane, using the covariant curvature tensor formalism developed in \\cite{SMS}. We have found that if there is only one positive-tension de Sitter brane, the Einstein gravity is recovered on the brane without any corrections other than those from the Kaluza-Klein modes. There is no trace of the radion mode, at least at the linear order, contrary to our naive expectation that the displacement of the brane does have physical meaning for the de Sitter 3-sphere brane and hence that some trace of it should appear on the brane. If we recall the case of a vacuum bubble in the 4-dimensional spacetime\\cite{GarVil}, we find exactly the same phenomenon there; an observer on the wall cannot detect fluctuations of the wall, though the brane does fluctuate in the 4-dimensional sense.  To investigate the physics behind this phenomenon further, as well as to clarify the role of the scalar field $\\phi$, we have derived the action for the displacement of the brane, $\\varphi$, coupled to the metric perturbation, $h_{\\mu\\nu}$, in Appendix A. The equation of motion of $\\varphi$ is then found to be exactly the same as that of $\\phi$ in the Randall-Sundrum gauge. This result may be interpreted as follows. The brane does fluctuate in this gauge by the amount $\\varphi(x)$ in the $(n+1)$-dimensional sense. However, the displacement of the brane cannot be detected by itself because one can always choose a local coordinate system in which $\\varphi=0$. {}From the view point of an observer on the brane, the natural coordinate associated with it is the gaussian normal coordinates defined by $h_{ab}n^b=0$ and $\\varphi=0$. We have then found $\\phi$ describes the intrinsic $n$-dimensional scalar curvature perturbation of the brane. Hence the effect due to the fluctuation of the brane cannot be detected unless it causes a non-trivial metric perturbation in the bulk and propagates back to the brane (or to the other brane in the two-brane case). Since the physical metric perturbation in the bulk is described by $E_{\\mu\\nu}$, the fact that the Green function in the zero-mode truncation, Eq.~(\\ref{zerotrunc}), vanishes in the limit of a single brane implies that the bulk metric becomes increasingly rigid as the location of the negative-tension brane approaches $z_{-}\\to\\infty$ and the fluctuation ceases to propagate through the bulk. As a result, the gravity on the positive-tension brane is not affected by the fluctuation of the brane, $E_{\\mu\\nu}=0$, in the single brane case.  In the two-brane system, on the other hand, there exists the cosmological radion mode as in the flat two-brane system. An interesting fact is that the cosmological radion has the negative mass-squared $m^2=-n{\\cal K}$. This may have important implications to the cosmological perturbation theory. In connection with this issue, we should mention that our result that the radion has the negative mass-squared is in apparent contradiction with the result obtained by Chiba\\cite{Chiba}, who derived a non-linear effective action for the radion mode by assuming a certain form of the 5-dimensional metric. Naively applying his result to curved background and taking the large separation limit, one finds the radion is a massless conformal scalar, i.e., it effectively has a positive mass-squared on the de Sitter brane. However, it should be noted that the assumed form of the 5-dimensional action used in \\cite{Chiba} is valid only for the flat brane background. So, there is in fact no logical contradiction. To understand the role of the radion in the inflationary brane universe, it will be necessary to derive the effective action of the radion on the de Sitter brane. This issue is currently under investigation."
    },
    "0011/astro-ph0011515_arXiv.txt": {
        "abstract": "We review recent observations of circularly polarized radiation from AGN made with the VLBA and with the ATCA. We also discuss briefly the detections of the Galactic sources Sag A* and SS433. The origin of the circular polarization is still an open question in most cases, and we discuss four possible mechanisms. Detectable circular polarization is a common property of quasars, but not of radio galaxies, and is always associated with the compact core. There is growing evidence that the {\\em sign} of the circular polarization stays the same over at least two or three decades in time, suggesting it is a fundamental property of the jet. ",
        "introduction": "Circular polarization (CP) measurements of extragalactic radio sources have long enjoyed the dubious reputation of being challenging to carry out and hard to interpret. After a period of activity in the 1970's and early 80's (summarized by Weiler \\& de Pater 1983) there have been few new observations until recently.  The renewed interest stems from two developments. First, the Australia Telescope Compact array (ATCA) can now measure CP with a limiting accuracy as low as $0.01\\%$ (Rayner, Norris \\& Sault 2000). This remarkable precision is because the array was specifically designed to measure CP, and uses orthogonal {\\em linearly} polarized feeds at each antenna. Second, the excellent performance of the VLBA has permitted the detection of CP in more than 20 AGN (Wardle et al. 1998, Homan \\& Wardle 1999, Homan, Attridge \\& Wardle in preparation). Because the VLBA is equipped with {\\em circularly} polarized feeds, the limiting accuracy is $\\sim 0.1\\%$, some ten times worse than ATCA. In compensation, the VLBA has about one thousand times higher resolution. Most measurements have been of AGN, but we note that CP has also been detected in the Galactic sources Sag A$^*$ using the VLA (Bower, Falke \\& Backer 1999) and SS433 using ATCA (Fender et al. 2000).  The promise of CP measurements is that they may permit, depending on the emission mechanism, direct measurement of the magnetic field strength and its polarity, the true magnetic flux carried by a jet (and hence generated at the central engine), the low energy cut-off in the relativistic electron spectrum, and the composition of the radiating plasma - whether it is predominantly an electron-proton plasma or an electron-positron plasma, or a mixture of the two. None of these can be determined directly from measurements of the first three Stokes parameters (total intensity, I, and linear polarization, Q and U). The fourth Stokes parameter, V, corresponding to the circularly polarized component of the radiation field, therefore offers a new way of measuring fundamental physical properties of AGN, and also the smaller scale jets in Galactic X-ray binaries. ",
        "conclusions": "It was stated in the introduction that ``circular polarization (CP) measurements of extragalactic radio sources have long enjoyed the dubious reputation of being challenging to carry out and hard to interpret.'' This is still true. It is hoped that this review will convince some readers that the unique information contained in such observations makes them, nevertheless, extremely valuable and worthwhile."
    },
    "0011/hep-ph0011399_arXiv.txt": {
        "abstract": "A degenerate neutron gas in equilibrium with a background of electrons and protons in a magnetic field exerts its pressure anisotropically, having a smaller value perpendicular than along the magnetic field. For critical fields the magnetic pressure may produce the vanishing of the equatorial pressure of the neutron gas, and the outcome could be a transverse collapse of the star. This fixes a limit to the fields to be observable in stable pulsars as a function of their density. The final structure left over after the implosion might be a mixed phase of nucleons and meson ($\\pi^{\\pm,0},\\kappa^{\\pm,0}$) condensate (a strange star also likely) or a black string, but no magnetar at all. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/gr-qc0011002_arXiv.txt": {
        "abstract": "We present the concept of a sensitive {\\itshape{and}} broadband resonant mass gravitational wave detector. A massive sphere is suspended inside a second hollow one. Short, high-finesse Fabry-Perot optical cavities read out the differential displacements of the two spheres as their quadrupole modes are excited. At cryogenic temperatures one approaches the Standard Quantum Limit for broadband operation with reasonable choices for the cavity finesses and the intracavity light power. A molybdenum detector of overall size of 2 m, would reach spectral strain sensitivities of $2\\cdot 10^{-23}$~Hz$^{-1/2}$ between 1000~Hz and 3000~Hz. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011045_arXiv.txt": {
        "abstract": "Astronomers are always anxious to push their observations to the limit -- basing results on objects at the detection threshold, spectral features barely stronger than the noise, or photometry in very crowded regions.  In this paper we present a careful analysis of photometry in crowded regions, and show how image blending affects the results and interpretation of such data.  Although this analysis is specifically for our NICMOS observations in M31, the techniques we develop can be applied to any imaging data taken in crowded fields; we show how the effects of image blending will even limit NGST.  We have obtained HST-NICMOS observations of five of M31's most metal rich globular clusters.  These data allow photometry of individual stars in the clusters and their surrounding fields.  However, to achieve our goals -- obtain accurate luminosity functions to compare with their Galactic counterparts, determine metallicities from the slope of the giant branch, identify long period variables, and estimate ages from the AGB tip luminosity, we must be able to disentangle the true properties of the population from the observational effects associated with measurements made in very crowded fields.  We thus use three different techniques to analyze the effects of crowding on our data, including the insertion of artificial stars (traditional completeness tests) and the creation of completely artificial clusters.  These computer simulations have proven invaluable in interpreting our data.  They are used to derive threshold- and critical-blending radii for each cluster, which determine how close to the cluster center reliable photometry can be achieved.  The simulations also allow us to quantify and correct for the effects of blending on the slope and width of the RGB at different surface brightness levels.  We then use these results to estimate the limits blending will place on future space-based observations. ",
        "introduction": "\\label{sec:introduction}  The main objective of our observations was to obtain physical parameters for a selection of metal-rich globular clusters in M31.  These parameters are usually straight-forward to derive from color-magnitude diagrams (CMDs), however, we have found the effects of crowding to be particularly severe.  We therefore present our analysis in two parts. This paper describes in detail the effects of crowding, how we quantify them, the techniques used to correct for them, and their implications for future space-based observations such as with NGST.  A second paper \\cite[][hereafter Paper II]{SFFJ2001} presents the science obtained from our observations using these techniques.  The general effects of crowding are well known, although not always accounted for.  These effects include reduced photometric accuracy which broadens the CMD and smears the luminosity function (LF), reduced positional accuracy \\citep{Hog2000}, artificial brightening and shifts in the measured colors \\citep{AG1995}, and in severe cases, the creation of objects brighter than anything in the parent population through random groupings of many stars \\citep{Ren1998, DTFA1993}.  If these effects are not taken into account, they can result in a mistaken spread in metallicity assumed from the RGB width, an incorrect mass function extrapolated from the LF, or a false age estimation from the AGB peak luminosity.  This paper is organized as follows.  Section \\ref{sec:observations} presents the details of our observations, and section \\ref{sec:data_reduction} explains the reduction procedures.  Section \\ref{sec:artificial_star_tests} describes the different simulations in detail.  Section \\ref{sec:results} presents the results which are applied to the cluster data in Paper II, and section \\ref{sec:implications} shows how these results can be applied to future space-based observations, such as the NGST.  We summarize our main conclusions in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In order to interpret out HST-NICMOS observations of globular clusters in M31, we have developed techniques for understanding and quantifying the effects of blending on stellar photometry.  Traditional completeness tests, specifically the injection and recovery of a handful of artificial stars into a preexisting frame, are fine for estimating the photometric completeness.  Unfortunately, this type of test does not help in understanding the true brightnesses of stars seen in crowded regions.  In very crowded regions, the few added artificial stars are merely added on top of preexisting clumps, with no way of recovering the composition of those clumps.  Our solution is to simulate the entire frame.  Our cluster simulations, consisting of hundreds of thousands of stars each (Fig. \\ref{fig:a280}), revealed that stars as bright as we observe near the cluster centers, can be easily created by the blending of many fainter stars (Fig. \\ref{fig:a280cmd}).  These simulations can be tailored to nearly any observation, and use any desired input LF to best determine the true properties of the observed stellar population, although where blending is severe, there are degeneracies in the shape of the input LF.  Using uniformly populated fields of varying surface brightness, we measured the degradation of photometric accuracy as a function of SB. We determined that the photometry of faint stars ($M_K \\sim -3$) began to be affected by blending at a surface brightness of $\\mu_K \\sim 16$ magnitudes arcsecond$^{-2}$.  At a surface brightness of $\\mu_K \\sim 14$ the photometry of even the brightest stars was affected by blending (Fig. \\ref{fig:artfieldsb}).  With these criteria, we determined threshold- and critical-blending radii for each cluster, which determine the proximity to each cluster where reliable photometry can be achieved (Tab. \\ref{tab:blend_radii}).  We show the effects of blending on the measured luminosity function.  At low surface brightnesses the measured LF accurately reflects the input LF.  However with increasing SB, the measured LF becomes skewed toward the bright end.  At very high SBs blending dominates the photometry, and the measured LF has a slope opposite in sign to the input LF, actually increasing toward bright stars (Fig. \\ref{fig:artfieldlfs}).  Blending also affects the measured slope and width of the giant branch. We quantified these effects, and determined a relation between the field SB and the artificial increase in the metallicity as derived from the RGB slope (Fig. \\ref{fig:gbslope}).  Results of this modeling will be applied to the cluster and field observations in Paper II.  By scaling our limiting surface brightness imposed by blending, we estimate the limitations blending places on future space-based observations.  We show that infrared observations of Cepheid variables with NGST will be trivial in the Virgo cluster, but may be nearly impossible in the Coma cluster.  \\"
    },
    "0011/astro-ph0011273_arXiv.txt": {
        "abstract": "A comparative study of the properties of the anelastic and subseismic approximations is presented. The anelastic approximation is commonly used in astrophysics in compressible convection studies whereas the subseismic approximation comes from geophysics where it is used to describe long-period seismic oscillations propagating in the Earth's outer fluid core. Both approximations aim at filtering out the acoustic waves while retaining the density variations of the equilibrium configuration. However, they disagree on the form of the equation of mass conservation. We show here that the anelastic approximation is in fact the only consistent approximation as far as stellar low-frequency oscillations are concerned.  We also show that this approximation implies Cowling's approximation which neglects perturbations of the gravity field. Examples are given to assert the efficiency of the anelastic approximation. ",
        "introduction": "Initiated with the work on atmospheric convection by Ogura \\& Phillips (1962) and Gough (1969), the anelastic approximation has been widely used in astrophysics to describe compressible convection.  For example, it has been first applied by Latour et al. (1976) and Toomre et al. (1976) to investigate the non-rotating and non-magnetized convective envelopes. Then, Gilman \\& Glatzmaier (1981) and Glatzmaier \\& Gilman (1981a,b,c) extended these studies to the rotating case whereas the MHD case has been considered by Glatzmaier (1984,1985a,b), Lantz \\& Sudan (1995) and Lantz \\& Fan (1999).  In all these studies, this approximation has been preferred to an other famous one in fluid dynamics, the so-called Boussinesq approximation where the velocity is assumed to be solenoidal, i.e. $\\na\\cdot \\vi = 0$ (Spiegel \\& Veronis 1960). As the density varies on several orders of magnitude in a convection zone, using the anelastic approximation is indeed more relevant and leads to $\\na\\cdot (\\rho \\vi) = 0$. However, both approximations filter out short-period acoustic oscillations by assuming that the Mach number of the convection is small so that `the task of numerical solution will not be complicated by the need to resolve very rapid time variations' (Latour et al. 1976).  More recently, the anelastic approximation has been used in the field of stellar oscillations. Waves can propagate over long distances in star interiors and large variations of density need also to be taken into account. Thus, as in stellar convection studies, the Boussinesq approximation is too restrictive and not adapted for this problem. However, since acoustic waves are filtered out, only the low-frequency part of the oscillations spectrum is captured. As this part of the spectrum also corresponds to the most perturbed one when rotation acts, this approximation is very attractive when low-frequency modes of a rotating star are considered.  Understanding the spectrum of rotating stars is indeed a difficult problem when the rotation period is of the same order than the oscillation one. In this case, the usual perturbative theory of Ledoux (1951) reaches its limits and the eigenvalue problem needs to be solved non-perturbatively (Dintrans \\& Rieutord 2000, hereafter referred to as Paper I). But rotation introduces new computational challenges: ({\\em i}) the strong Coriolis coupling between normal modes of different degrees $(\\ell,\\ell\\pm 1)$ leads to a large system of coupled differential equations; ($ii$) some singularities due to the presence of wave attractors emerge in the region $\\sigma \\simeq 2 \\Omega$, $\\sigma$ being the wave frequency and $\\Omega$ the uniform rotation rate of the star. Applying the anelastic approximation does not remove these singularities but decreases the size of the numerical problem since acoustic quantities disappear (see Paper I). We note that Berthomieu et al. (1978) used the same arguments when considering a hierarchy between the displacement and pressure components, this hierarchy leading to the same result as the direct use of the anelastic approximation.  Contrary to stars, the mass of the Earth is not strongly concentrated near its centre. Only 30\\% of the whole mass is indeed inside the half radius whereas this ratio is around 90\\% for the Sun (see e.g. Stix 1989). It means that the variations in density in the Earth's outer fluid core generated by the propagation of seismic P-waves may produce non negligible variations in the gravitational potential $\\phi$ and explains why geophysicists do not wish to use the Cowling approximation in their description of P-waves. To take into account this effect but to simplify the equations for low frequency modes Smylie \\& Rochester (1981), Smylie, Szeto \\& Rochester (1984) introduced the `subseismic' approximation where self-gravity is included.  Recently, De Boek, Van Hoolst \\& Smeyers (1992) applied this approximation to describe the low-frequency g-modes of non-rotating stars. Starting from the same equations as Smylie \\& Rochester (1981), they found an analytic expression for the {\\em asymptotic} low frequencies which slightly differs from the standard one deduced by Tassoul (1980) in the Cowling approximation. By applying this expression to the low-frequency oscillations of both homogeneous and polytropic ($n=3$) star models, De Boek et al. (1992) concluded that the subseismic approximation is more accurate than the standard one in the polytropic case but less accurate in the homogeneous case.  However, as will be shown below, the subseismic approximation is not a consistent approximation of the equations of motion; we will show indeed that the proper way to simplify the equations in order to deal with the low frequency modes is to use the anelastic approximation and that this approximation also implies Cowling's approximation.  Hence, after establishing the complete set of equations describing linear oscillations of a stratified compressible fluid and the appropriate boundary conditions (Section 2), we focus on the similarities and differences of the anelastic and subseismic approximations (Section 3); in particular, we show that both imply Cowling's approximation and demonstrate the inconsistency of the subseismic approximation. Then we compare the two approximations on the test cases used by De Boek et al. (1992), namely, the homogeneous and polytropic star models (Section 4). Comparison of the resulting approximate eigenfrequencies with the exact ones (computed numerically for the polytrope but analytically for the homogeneous model) shows clearly the better efficiency of the anelastic approximation. Finally, we conclude in Section 5 with some outlooks of our results. ",
        "conclusions": "We studied the properties of the subseismic and anelastic approximations devised for low-frequency g-modes of stars.  These approximations both assume that the Eulerian perturbations of pressure do not contribute to their Lagrangian parts. As a consequence, acoustic waves are filtered out. However, the equation of mass conservation differs since the subseismic approximation keeps the contribution coming from the Eulerian density fluctuation. We showed that this is incorrect because it is inconsistent with the neglect of the Eulerian pressure perturbation.  As an illustration we first applied both approximations to the low-frequency g$^-$-modes of the homogeneous star. In this case, we found an analytic expression for the subseismic and anelastic eigenfrequencies using power-series solutions.  We thus showed that the anelastic approximation is more precise than the subseismic one as expected.  As a second test, we considered the low-frequency g$^+$-modes of the polytrope $n=3$. No analytic solutions can be found for this model and eigenfrequencies have been computed numerically using an iterative solver.  We thus computed the subseismic and anelastic eigenfrequencies and compared them to the exact ones. As in the homogeneous case, the anelastic approximation is more accurate than the subseismic one.  We also showed that Cowling's approximation is a necessary consequence of the anelastic approximation; it is therefore unnecessary to enforce this approximation since perturbations of self-gravity will automically decouple from the set of equations and will not influence the eigenfrequencies.  As already mentioned in the introduction, an obvious application of the anelastic approximation concerns the low-frequency oscillations of rapidly rotating stars. Applying this approximation indeed removes the acoustic quantities and dramatically decreases the size of the numerical problem. An example is given in Paper I with the study of the oscillations of a rapidly rotating $\\gamma$ Doradus-type star, for which periods of oscillations are of the same order as that of rotation (i.e. around one day).  \\begin{figure} \\psfig{file=ma965_fig2.ps,height=7cm} \\caption{Rates of relative errors of computed eigenfrequencies for the polytrope $n=3$.} \\label{fig2} \\end{figure}"
    },
    "0011/astro-ph0011429_arXiv.txt": {
        "abstract": "We examine the power spectra of IDV and show the information, which is to be gained by wavelet analysis of light curves of the quasar 0917+624. Results for total and polarized flux at 11\\,cm are shown. Both interstellar scattering and intrinsic models have difficulties in explaining the 1 day period variations. A theoretical model for the time averaged emission is presented, which provides the basis for the analysis of possible variations. ",
        "introduction": "Intraday variability (IDV) in extragalactic radio sources has been found in sources like 0917+624 (Quirrenbach et al. 1989) with variations on time scales of about 1\\,day (fast IDV), showing amplitudes of 5-25\\% in total flux. The brightness temperatures of these sources range from $10^{18}$\\,K to $10^{21}$\\,K, if the variability time scale measures the source size. More recently ultrafast IDV sources PKS 0405-385 (Kedziora-Chudczer et al. 1997) and J1819+3845 (Dennett-Thorpe \\& de Bruyn 2000) showing time scales of half an hour to a few hours and much larger amplitudes have been found. While refractive interstellar scintillation (RISS) seems to be the only successful explanation for ultrafast IDV sources, the problem has not been solved for longer period IDVs.  We compared the results of Fourier power spectra and a wavelet analysis of light curves. Published data are used (e.g. Qian et al. 1991) from a 1989 multi-frequency campaign with the VLA and the 100\\,m Effelsberg antenna at 3.6, 6, 11 and 20\\,cm of total and polarized flux. The quasar 0917+624 has a determined redshift of 1.44 and shows superluminal motion up to\\footnote{We assume a matter dominated universe with Hubble constant $H_0 = 100h$ km s$^{-1}$ Mpc$^{-1}$ and deceleration $q_0 = 0.5$.} $h\\beta_{app} \\la 8$ (Standke et al. 1996). ",
        "conclusions": ""
    },
    "0011/astro-ph0011335_arXiv.txt": {
        "abstract": "The {\\it backward evolution} approach to modelling galaxy source counts is re-visited in the wake of the numerous results and revelations from the Infrared Space Observatory (ISO), the Submillimetre Common User Bolometer Array (SCUBA) and the detections and measurements of the cosmic extragalactic background light. Using the framework of the Pearson \\& Rowan-Robinson (1996) galaxy evolution model, the observed source counts and background measurements are used to constrain the evolution in the galaxy population. It is found that a strong evolution in both density and luminosity of the high luminosity tail of the IR luminosity function, interpreted as the ultraluminous galaxies discovered first by IRAS and later elevated in status by SCUBA and ISO, can account for the source counts from 15$\\umu$m (where it matches the undulations in the integral counts and hump in the differential counts extremely well), to the sub-mm region, as well as explain the peak in the cosmic infrared background at $\\sim$140$\\umu$m. The sub-mm counts are interpreted as the superposition of 2 separate populations comprising of ULIGs at the brighter sub-mm fluxes and starburst galaxies at fluxes fainter than $\\sim$2mJy. In this scenario the high redshift ULIGs are tenuously interpreted as the progenitors of today's giant elliptical galaxies.  All the source count models can be accessed via the world wide web at this URL :  $http://www.ir.isas.ac.jp/\\sim cpp/counts/$ ",
        "introduction": "\\label{sec:introduction}  The formation epoch of galaxies has long been the {\\it holy grail } of cosmology. Traditionally one of the most popular and successful ways of glimpsing this epoch has been the analysis of extragalactic source counts or more commonly galaxy number counts. Although originally intended as a means of determining the geometry of the Universe this tool has shown to be much more viable and useful in the study of the evolutionary and star formation history of galaxies (Kirshner et al.\\shortcite{kirshner81}, Ellis\\shortcite{ellis87}). In particular, the large dust content of galaxies discovered by IRAS and the strong K-corrections have made the infra-red and more recently, the sub-mm, extremely important wavebands for the study of galaxy evolution and cosmic star formation. By modelling the counts of galaxies it has become possible to analyze the evolution of the galaxy population out to redshifts of unity and higher. Coupled with the advent of new observational data, galaxy number counts are poised to discriminate between rivalling evolutionary theories.  Recently, it has become the vogue to tackle the problem of the number distribution of galaxies in the Universe via one of two methods. The first (often called the {\\it backward evolution} approach \\cite{lonsdale96}), takes the observed, present day ($z=0$) luminosity function and evolves it in luminosity and/or density back out to higher redshifts assuming some parameterization of the evolution (Beichman \\& Helou \\shortcite{beich91}, Blain \\& Longair \\shortcite{blain93}, Pearson \\& Rowan-Robinson \\shortcite{cpp96}, Xu et al. \\shortcite{xu98}, Takeuchi et al. \\shortcite{take99}). This method has the advantages that it is both direct and relatively simple to implement with few free parameters and assumptions about the Universe at earlier times. The disadvantage in the past was that the information on which the assumptions about evolution were made often came from IRAS data which extended out only to relatively low redshifts $\\approx 0.02-0.2$. In the alternative method, known as the {\\it forward evolution} approach the evolution is computed assuming some initial conditions of the physical processes of chemical evolution and photometric evolution of the stellar populations that heat the dust \\cite{fran94}. The so-called {\\it semi analytical approach} combines these assumptions on the chemical/photometric evolution of galaxies with models considering the dissipative and non-dissipative processes affecting galaxy formation within dark matter haloes (Cole et al. \\shortcite{cole94}, Baugh et al \\shortcite{baugh96}) and has provided a reasonable fit to the source counts in the infra-red \\cite{guid98}. The {\\it forward evolution} approach has the advantage of being based on perhaps a more fundamental set of assumptions, however the obvious disadvantages are the larger number of free/unknown parameters assumed in these models.  Furthermore, with the recent advent of large, deep surveys by SCUBA (Holland et al. \\shortcite{holland99}, e.g.  Smail et al.\\shortcite{smail97}, Hughes et al.\\shortcite{hugh98}, Barger et al.\\shortcite{barg98}, Blain et al.\\shortcite{blain99}, Barger et al.\\shortcite{barg99}) and ISO (Kessler et al.\\shortcite{kessler96}, e.g. Bogun et al.~\\shortcite{bogun96}, Serjeant et al.~\\shortcite{serjeant97}, Taniguchi et al.~\\shortcite{taniguchi97}, Elbaz et al.~\\shortcite{elbaz98}, Stickel et al.~\\shortcite{stick98}, Kawara et al.~\\shortcite{kawara98}, Flores et al.~\\shortcite{flores99a},~\\shortcite{flores99b}, Altieri et al.~\\shortcite{altieri99}, Oliver et al.~\\shortcite{oliver00a}, Lindern-V{\\o}rnle et al.~\\shortcite{vornle00}, Oliver et al.~\\shortcite{oliver00b}), including constraints on the source counts to fainter levels and the IR background (Lagache et al.\\shortcite{lagache99},\\shortcite{lagache00}, Matsuhara et al. \\shortcite{mat00}, Matsumoto et al.\\shortcite{mats00}) and detections by the COBE FIRAS/DIRBE instruments at 140 \\& 240$\\umu$m (Puget et al. \\shortcite{pug96}, Fixsen et al. \\shortcite{fixsen98}, Hauser et al. \\shortcite{hauser98}) the tools are now available to constrain the {\\it backward evolution} methodology to significantly higher redshifts. In the light of these recent advances in the mm-IR region the {\\it backward evolution} method is re-visited.  The aim of this paper is not to produce a definitive evolutionary model to fit all observed counts while satisfying various preferences on the various dominant populations (starburst, AGN, luminous infrared galaxies (LIG's $L_{IR} > 10^{11}L_{\\sun}$), ultraluminous galaxies (ULIG's $L_{IR} > 10^{12}L_{\\sun}$), hyperluminous galaxies ($L_{IR} > 10^{13}L_{\\sun}$)). Rather I seek to investigate the various avenues of evolution that can viably fit the galaxy source counts from the mid-IR to sub-mm wavelengths whilst not violating the constraints set by the IR background and CMB measurements. The paper is structured as follows, in section \\ref{sec:model} the galaxy count and world model parameters are outlined. Section \\ref{sec:evol} reviews current evolutionary models and constraints and analyszes simple density and luminosity evolution scenarios. Section \\ref{sec:counts} formulates a new evolutionary source count model and applies it to the observed source counts and background measurements from the mid-IR to sub-mm wavelengths. The possible interpretations and conclusions are reported in sections \\ref{sec:discuss} \\& \\ref{sec:conc} respectively. ",
        "conclusions": "\\label{sec:conc}  The {\\it backward evolution} approach to galaxy evolution has been re-visited in light of the recent observations by ISO and SCUBA in the infra-red and sub-mm regions. Assuming classical evolutionary scenarios incorporating power law luminosity and density evolution it is impossible to simultaneously fit all the source counts from 15-850$\\umu$m and the CIRB.  A new evolutionary model is investigated which assumes the foundations of the Pearson \\& Rowan-Robinson ~\\shortcite{cpp96} model (cppRR). In addition to this original framework, a strongly evolving high luminosity IR galaxy component (referred to as the ULIG component but specifically the high luminosity end of the 60$\\umu$m luminosity function) is added. This new population evolves in both density and luminosity and is capable of fitting both the source counts from 15-850$\\umu$m and the cosmic far-infrared background radiation, including the peak at $\\approx$140$\\umu$m and the long wavelength tail.  In this scenario the ULIG component is tenuously interpreted as the progenitor of the present day giant elliptical galaxy population.  The excesses in the 15$\\umu$m and 170$\\umu$m counts are well explained by strong evolution to redshift $\\sim$1 of a population of luminous, $lg L_{IR} \\sim 12 L_{\\sun}$, infrared galaxies. This population of galaxies becomes prominent at flux levels just below the limits of the IRAS 60$\\umu$m survey (although previously hinted at) and can be associated with the archetypical ULIG Arp220. These results are generally consistent with the numerous surveys at 15$\\umu$m ~\\cite{elbaz00} which found that the dominant population above a redshift of $\\sim$0.5 were massive ($M>10^{11}M_{\\sun}$, luminous $L_{IR} > 10^{11-12}L_{\\sun}$ galaxies.  At 850$\\umu$m, these galaxies dominate the bright SCUBA counts at S$\\geq$6mJy. Towards fainter fluxes the starburst galaxies begin to contribute strongly to the source counts, becoming dominant at fluxes $<2mJy$. In this scenario the sub-mm counts can be divided into 2 distinct populations, the dusty, giant ULIGs which would be considered the progenitors of giant elliptical galaxies and the IRAS starburst population / Lyman break galaxies. This conclusion is in good agreement with the analysis of Busswell \\& Shanks~\\shortcite{busswell00} who transformed the optical counts of faint starburst galaxies to the sub-mm waveband. Using a Bruzual \\& Charlot \\cite{bruzchar93} galaxy model with an exponentially decaying SFR of $\\tau = 9Gyr$, their models successfully predicted optical results, including the Lyman break galaxies \\cite{metcalfe00}. Furthermore they accounted for the sub-mm counts below a sensitivity of $\\sim$1mJy but failed to fit the counts at brighter fluxes, concluding that an additional population was required (ULIGs/AGN?). Of further interest is that their models were also deficient in fitting the CIRB at 140-240$\\umu$m where the ULIG population is most prominent.  Note, analysis of the so called {\\it Madau} plot of the star formation history of the galaxies in this evolutionary model will be addressed in a later paper. The primary aim of this work is to ascertain what evolutionary scenarios can succesfully fit the observed source counts / background radiation.  From an analysis of luminosity and density evolution the following conclusions can be drawn from the galaxy source counts.  \\begin{enumerate} \\item The mid-IR region (15$\\umu$m) is sensitive to evolution in the low redshift population but is not strongly constrained by evolution at higher redshift (z$>$1) \\item The far-IR region at 60-90$\\umu$m does not presently include enough information to constrain the evolution in the source counts. \\item The far-IR region at 170$\\umu$m is sensitive to the low redshift population and the faint end counts constrain the contribution to the peak of the CIRB. \\item The sub-mm counts constrain the magnitude of the evolution in the high redshift Universe z$>$1 but due to the steep K-corrections, cannot distinguish evolutionary forms between redshifts 1-2. \\end{enumerate}  Using the short wavelength 15$\\umu$m and long wavelength 850$\\umu$m counts, together with the CIRB as bounding conditions, the evolution over a wide range of both wavelength and redshift can be constrained. It is found that a scenario where the ULIG population undergoes 2 phases of evolution, a rapid merging epoch to z$\\sim$1 and exponential evolution in luminosity to higher redshifts, best fits the observations.  In this scenario the high redshift ULIGs detected at 850$\\umu$m fluxes $\\geq$ 2mJy may be interpreted as the progenitors of todays giant elliptical galaxies while the population at fluxes fainter than $S_{850} \\leq$2mJy represent the IRAS starburst population / Lyman break galaxies.  Observations in the next few years in the mid-IR with SIRTF \\cite{rieke00} and the Japanese ASTRO-F mission (Matsuhara et al.~\\shortcite{mat98}, Watarai et al. ~\\shortcite{wat00}, Pearson et al.~\\shortcite{cpp01}), in the far-IR-sub-mm-mm with FIRST \\cite{pilbrat00}, LMT \\cite{schloerb97} and an all sky survey at far-IR wavelengths with ASTRO-F (Kawada et al.~\\shortcite{kaw98}, Takahashi et al. \\shortcite{takahashi00}) will undoubtebly solve many of these mysteries."
    },
    "0011/astro-ph0011321_arXiv.txt": {
        "abstract": "We perform direct numerical simulations of three-dimensional freely decaying magnetohydrodynamic (MHD) turbulence. For helical magnetic fields an inverse cascade effect is observed in which power is transfered from smaller scales to larger scales. The magnetic field reaches a scaling regime with self-similar evolution, and power law behavior at high wavenumbers. We also find power law decay in the magnetic and kinematic energies, and power law growth in the characteristic length scale of the magnetic field. ",
        "introduction": "Within cosmology, astrophysics or geophysics one often needs to deal with electrically conducting plasmas at high kinematic and magnetic Reynolds numbers where magnetic fields are dynamically important. Indeed, much of the turbulence in the interstellar medium is magnetohydrodynamic in nature.  Hydromagnetic turbulence has been explored extensively in connection with the generation of large scale magnetic fields in astrophysical bodies such as planets, stars, accretion discs and galaxies through dynamo theories. Non-driven, freely decaying turbulence may also be of interest in connection with both the physics of the interstellar medium and cosmology. Our interest was inspired by the cosmology of primordial magnetic fields, which is sometimes considered as a possible source for providing the seed for the galactic dynamo \\cite{Zel88}.  There have been various related works on decaying MHD turbulence, by authors interested in different contexts \\cite{Hos+95,PolPouSul95,GalPolPou97,Mac+98,BisMul99,MulBis00}. Most directly comparable to our work, Biskamp and M\\\"uller \\cite{BisMul99} studied the energy decay in incompressible 3D magnetohydrodynamic turbulence in numerical simulations at relatively high Reynolds number, and in a companion letter \\cite{MulBis00} studied the scaling properties of the energy power spectrum.  We are here especially interested in the inverse cascade of magnetic helicity whereby magnetic energy is transferred from small to large scale fluctuations. This is important for a primordial magnetic field to reach a large enough scale with sufficient amplitude to be relevant for seeding the galactic dynamo \\cite{HinEve98}.  It should be noted that due to the conformal invariance of MHD in the radiation era the MHD equations in an expanding universe can be converted into the relativistic MHD equations in flat spacetime by an appropriate scaling of the variables and by using conformal time \\cite{BraEnqOle96}. The equations of \\cite{BraEnqOle96} differ slightly from the ordinary non-relativistic MHD equations. However, in order to facilitate comparison with earlier work, we use the non-relativistic equations.  We perform 3D simulations both with and without magnetic helicity, starting from statistically homogeneous and isotropic random initial conditions, with power spectra suggested by cosmological applications. We find a strong inverse cascade in the helical case, with equivocal evidence for a weak inverse cascade when only helicity fluctuations are present. In the helical case we also find a self-similar power spectrum with an approximately $k^{-2.5}$ behavior at high $k$. We present energy decay laws which are comparable to those found in the incompressible case by Biskamp and M\\\"uller \\cite{BisMul99}, and in the compressible case by Mac Low et al.\\ \\cite{Mac+98}. ",
        "conclusions": "We have shown that for an isothermal and compressible magnetized turbulent fluid, when undergoing a process of free decay, a substantial inverse cascade is present for helical magnetic field configurations, which transfer energy from smaller scale magnetic fluctuation to larger scale ones. For non-helical magnetic fields only a weak inverse cascade was observed on the largest scales.  The energy spectrum of the magnetic field shows evidence for a self-similar evolution with a development of a power law of roughly $k^{-2.5}$ beyond the peak. Decay laws for both the kinematic and magnetic energy were found. The kinetic energy decay was approximately $t^{-1.1}$ for both helical and non-helical magnetic fields. The decay of the magnetic field energy was found to be strongly dependent on the the initial helicity, decaying roughly as $t^{-0.7}$ and $t^{-1.1}$ for helical and non-helical initial conditions respectively. For the helical case, the magnetic energy decay rate showed a dependence on the Reynolds number, with a slower decay rate for larger Reynolds numbers.  We also observed power law behavior in the characteristic length scale of the magnetic field, defined as the Taylor microscale $L_{T}$. In the helical case $L_{T} \\sim t^{0.5}$, whereas for non-helical fields the growth was somewhat slower, $L_{T} \\sim t^{0.4}$, and we ascribe the faster growth rate to the presence of the inverse cascade in the helical case."
    },
    "0011/astro-ph0011117_arXiv.txt": {
        "abstract": "{ Results from observations of the young oxygen-rich supernova remnant \\object{SNR\\,0102$-$72.3} in the Small Magellanic Cloud during the calibration phase of the XMM-Newton Observatory are presented. Both EPIC-PN and MOS observations show a ringlike structure with a radius of $\\sim$15\\arcsec\\ already known from Einstein, ROSAT and Chandra observations. Spectra of the entire SNR as well as parts in the eastern half were analyzed confirming shocked hot plasma in non-uniform ionization stages as the origin of the X-ray emission. The spectra differ in the northeastern and the southeastern part of the X-ray ring, showing emission line features of different strength. The temperature in the northeastern part is significantly higher than in the southeast, reflected by the lines of higher ionization stages and the harder continuum. Comparison to radio data shows the forward shock of the blast wave dominating in the northern part of the SNR, while the southern emission is most likely produced by the recently formed reverse shock in the ejecta. In the case of the overall spectrum of \\object{SNR\\,0102$-$72.3}, the two-temperature non-equilibrium ionization model is more consistent with the data in comparison to the single plane-parallel shock model. The structure of \\object{SNR\\,0102$-$72.3} is complex due to variations in shock propagation leading to spatially differing X-ray spectra. ",
        "introduction": "\\object{SNR\\,0102$-$72.3} is one of the brightest X-ray sources in the Small Magellanic Cloud (SMC) and was discovered by the IPC instrument onboard the Einstein Observatory (Seward \\& Mitchell \\cite{SM81}). The Einstein HRI resolved a shell-like supernova remnant (SNR), and optical emission is seen from a variety of filaments rich in oxygen and neon (Dopita \\& Tuohy \\cite{DT84}) arranged in an incomplete clumpy ring with radius of about 12\\arcsec. Enrichment in the element oxygen is indicative for remnants of a type II SN explosion. The maximum velocity of the optical filaments [\\ion{O}{iii}], not visible in H$\\alpha$ (Dopita \\et\\ \\cite{D81}), was found to be up to 6500\\,km/s (Tuohy \\& Dopita \\cite{TD83}). The bright \\ion{O}{iii} knots, assumed to be dense ejecta clumps or Rayleigh-Taylor fingers of ejecta (density $\\sim$1\\,cm$^{-3}$), correspond in some parts well with bright X-ray features (Gaetz \\et\\ \\cite{G00}) but in other parts not. The estimated ionization timescale of $\\mathrm{log}(nt)<12$ adds to the origin of the X-ray emission being a shock heated X-ray plasma in a state of non-equilibrium ionization (NEI). Radio observations of \\object{SNR\\,0102$-$72.3} with 3\\arcsec\\ resolution in the 6\\,cm  wavelength band by Amy \\& Ball (\\cite{AB93}) unveiled an outer radio shell diameter of 40\\arcsec$\\pm$5\\arcsec. The radio emission is found to lie predominantly outside the bright X-ray emission but within an outer faint X-ray rim of $\\sim$44\\arcsec\\ (Gaetz \\et\\ \\cite{G00}). Hayashi \\et\\ (\\cite{H94}) reported on a broadband spectrum of \\object{SNR\\,0102$-$72.3} gained from the ASCA SIS instrument showing most of the emission from \\object{SNR\\,0102$-$72.3} lying in the 0.5 -- 2.3\\,keV energy band. The most prominent line features were interpreted as He-like K$\\alpha$, H-like K$\\alpha$ and  K$\\beta$ emission from ionized atoms of the elements oxygen, neon and  magnesium. Their difficulties to describe the overall spectrum with simple NEI models was interpreted as an inhomogeneous abundance pattern within the emitting plasma and possibly two components, namely forward shock in the interstellar medium (ISM) and reverse shock in the ejecta, contributing to the overall X-ray emission. XMM-Newton with its high sensitivity now enables spatially resolved spectroscopy on various parts of \\object{SNR\\,0102$-$72.3}. The following section describes the details of the observation and the data. Spectra from the selected regions are presented and fitted in Sec.\\,\\ref{xrayspec}. Finally, the results are discussed in Sec.\\,\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  \\object{SNR\\,0102$-$72.3} is a young SNR with an estimated age of about 1000\\,yr (Tuohy \\& Dopita \\cite{TD83}), which is no longer in the free expansion phase. Most of the X-ray emission originates from a bright ring with radius $\\sim$14\\arcsec\\ and a mean FWHM of 5\\arcsec\\ caused by a reverse shock propagating through the ejecta (Gaetz \\et\\ \\cite{G00}; Hughes \\et\\ \\cite{H00}). The XMM-Newton EPIC observations confirm this overall picture.  The spectra of the northeastern and southeastern regions of the SNR can be fitted with a two component NEI model, each component with a characteristic single ionization timescale. The emission lines of the two highest ionization stages of oxygen and neon, as well as the two values for $\\tau$ point out, that there is an ongoing shock ionization. This effect differs in the northeastern and the southeastern part of the SNR which can be seen in the unequal spectra of these regions (Fig.\\,\\ref{specej1} and Fig.\\,\\ref{specej2}). The plasma temperature for the northeastern part is up to 4 times higher than in the rest of the SNR, indicating that the shock velocity is higher in this region ($v_\\mathrm{s} \\sim T_\\mathrm{s}^{0.5}$). Though most of the X-ray emission of the SNR arises from the hot ejecta, in the northeastern part the forward shock of the blast wave propagating into the ISM becomes important, verified in radio observations showing synchrotron emission right behind the blast wave which can be seen in Fig.\\,\\ref{mosradio}. The X-ray emission of the entire SNR is dominated by the ejecta emission, which can be seen in the significant similarity between the overall spectrum and the spectrum of the southeastern region.  All the spectra extracted from EPIC-PN data were better reproduced by the two component NEI model than in a single plane-parallel shock model. The two NEI components differ not only in temperatures, but much more significantly in the ionization timescale $\\tau$. While the lower $\\tau_{2}$ component values themselves are similar for both selected regions as well as for the overall spectrum, the higher $\\tau_{1}$ value is the lowest in the northeastern region. In this part of the SNR, also the fractional emission measure of the higher $\\tau_{1}$ component is smaller than in the southeastern region or the whole SNR. Since in the southeastern (and the entire SNR) the ratio of the bright X-ray emission originating from the inner parts of the SNR (i.e.\\ the bright X-ray ring) to the X-rays from the outer parts is higher (see Fig.\\,\\ref{mosradio}), the higher $\\tau_{1}$ component can be assigned to the bright X-ray ring, outlining regions with higher densities. The implied density distribution around the SNR ring is also supported by the extended overall shape along the southwest to northeast axis, which was already reported by Gaetz \\et\\ (\\cite{G00}) and can be verified in Fig.\\,\\ref{mosradio}. This is indicative of a less decelerated expansion in that direction. Chandra observations have shown that there is evidence for a spatially varying distribution of the ionization stages and the existence of a reverse shock in the ejecta (Gaetz \\et\\ \\cite{G00}; Flanagan \\et\\ \\cite{F01}). The results obtained from the spectral analysis of the XMM-Newton EPIC-PN data show that there are differences in the plasma states between various parts of the SNR with spatial temperature and ionization stage variations, contributing to its complicated structure."
    },
    "0011/astro-ph0011267_arXiv.txt": {
        "abstract": "We simulate multi-color surveys, which use the same telescope time on different filter sets of broad-band and medium-band filters. We use a photometric classification method for identifying stars, galaxies and quasars and for estimating multi-color redshifts of extragalactic objects. A library of $> 65000$ color templates is used for comparison with observed objects. Spectroscopy from the CADIS survey has proven the performance of the method. We learned, that medium-band surveys perform superior to broad-band surveys in terms of classification and redshift estimation although they collect less photons. A full account of this work is already in print (Wolf, Meisenheimer, \\& R\\\"oser, 2001). ",
        "introduction": "The value of a multi-color survey is usually not simply limited by the depth of its imaging but by the depth of its classification used to select the objects of interest from the survey data. Here, we investigate what factors the depth of a successful classification depends on. Obviously, it depends in some way on the filter set used as well as on the depth of the imaging, but the optimum choices might vary according to the goal of the survey.  If a survey aims at identifying only one type of object with characteristic colors, a tailored filter set can be designed. E. g., when looking for U-band dropouts [9], the UGR filter set is a very good choice. The performance of such a dropout survey depends mostly on the depth reached in the U-band, so the photon flux detection limit in U is the key figure. Also, number count studies are limited by the completeness limit in the filter of concern.  If all kinds of different objects are the subject of the survey, a filter set covering the entire spectral range accessible to the survey instrument is a natural choice. The SLOAN Digital Sky Survey [14] is presently the most ambitious project to provide a broad-band {\\it ugriz} color database, on which the astronomical community might perform a large number of virtual surveys.  Of course, it is possible to split the bands of a filter set even further into a larger number of medium-band filters. Surveys like CADIS with typically 10 to 20 filters are sampling the visual spectrum with a resolution almost comparable to that of low resolution imaging spectroscopy. Intuitively, one might expect that such surveys are very expensive in terms of telescope time while their value might not be obvious. This paper will show, that in fact medium-band surveys are a powerful and economically competitive alternative to broad-band surveys. ",
        "conclusions": "We conclude, that medium-band surveys achieve at least the same working depth as broad-band surveys while providing much improved redshift resolution. They should provide high benefit already without consuming spectroscopic follow-up and be particularly well suited for multi-purpose surveys. Therefore, we recommend medium-band capabilities for future survey instruments and propose their use for public surveys.  Currently, we conduct three multi-color surveys, which make heavy use of medium-band filters. For each of them, filters and exposure times were chosen to match some primary survey strategy, which are not necessarily optimal choices in terms of a general classification. These are:  \\begin{enumerate} \\item The {\\it Calar Alto Deep Imaging Survey} (CADIS): Four broad-band (BRJK$^\\prime$) and twelve medium-band filters have mostly been chosen to match the needs of the emission line survey in CADIS, while some of them fill in gaps in the spectral coverage. The multi-color part of CADIS has been used to study the evolution of the galaxy luminosity function at $z=0.1\\ldots 1.3$, to search for quasars at all visible redshifts and to use the observed faint stellar population to check models of the Galactic structure and the stellar luminosity function (Meisenheimer et al. 1998, Wolf et al. 1999, Fried et al. 2000, Phleps et al. 2000, Wolf et al. 2000).  \\item A lensing study of the galaxy cluster Abell 1689: Two broad-band and seven medium-band filters have been chosen to separate well between the cluster galaxies at $z \\approx 0.19$ and the background population. The galaxy luminosity function in the background of the cluster is compared to a control field taken from CADIS, and the cluster mass is estimated from weak lensing effects on the apparent luminosities (Dye et al. 2000).  \\item A WFI survey for {\\it Classifying Objects by Medium-Band Observations with 17 filters} (COMBO-17): The filters (setup A from Tab.\\,1) are chosen to provide a selection function and a redshift accuracy, which is as independent of redshift as feasible. The data will be used to study the faint end of the quasar luminosity function at all accessible redshifts $z > 1$, galaxy-quasar correlation at $z<1.3$, as well as weak lensing effects in the cluster group Abell 901/2 and in the open galaxy field (Wolf et al. 2000). \\end{enumerate}"
    },
    "0011/astro-ph0011051_arXiv.txt": {
        "abstract": "In 1998 September, the X-ray transient XTE\\,J1550$-$564 underwent a major outburst in soft and hard X-rays, followed by a radio flare. Australian Long Baseline Array images obtained shortly after the peak in the radio flare showed evolving structure. The components observed have an apparent separation velocity of $>2c$. ",
        "introduction": "The X-ray transient XTE\\,J1550$-$564 was discovered simultaneously by the All-Sky Monitor on the Rossi X-ray Timing Explorer \\cite{smith} and by the Burst and Transient Source Experiment on the Compton Gamma-Ray Observatory \\cite{wilson} on MJD~51063 (1998~Sept~7; MJD$=$JD$-$240000.5). Following the discovery of the radio counterpart \\cite{cw}, XTE\\,J1550$-$564 was monitored with the Molonglo Observatory Synthesis Telescope (MOST) at 843~MHz between MJD~51065 and 51092. A flare was recorded, peaking with an estimated flux density of 380~mJy on MJD~51078. Two Very Long Baseline Interferometry (VLBI) observations using the Australian Long Baseline Array (LBA) were performed 2--3 days after the MOST flare. In this paper we present the resulting images from these observations, and discuss briefly the implications of the evolving structure observed. ",
        "conclusions": "We have presented two VLBI images of XTE\\,J1550$-$564 from the time of the 1998 September outburst and showed that the source evolved over a period of 24~hours. The two outermost components seen in the 1998 Sept 25 image are associated with the two components in the 1998 Sept 24 image, implying an apparent separation velocity of $>2c$."
    },
    "0011/astro-ph0011098_arXiv.txt": {
        "abstract": "We present the first interferometric study of the molecular gas in the metal--poor dwarf starburst galaxy NGC\\,4214. Our map of the $^{12}$CO (1--0) emission, obtained at the OVRO millimeter array, reveals an unexpected structural wealth.  We detected three regions of molecular emission in the north--west (NW), south--east (SE) and centre of NGC\\,4214 which are in very different and distinct evolutionary stages (total molecular mass: $\\sim 5.1 \\times 10^6 M_{\\odot}$).  These differences are apparent most dramatically when the CO morphologies are compared to optical ground based and HST imaging: massive star formation has not started yet in the NW region; the well--known starburst in the centre is the most evolved and star formation in the SE complex started more recently.  We derive a star formation efficiency of $\\sim8\\%$ for the SE complex. Using high--resolution VLA observations of neutral hydrogen (\\ion{H}{1}) and our CO data we generated a total gas column density map for NGC\\,4214 (\\ion{H}{1} + H$_2$). No clear correlation is seen between the peaks of \\ion{H}{1}, CO and the sites of ongoing star formation.  This emphasizes the irregular nature of dwarf galaxies.  The \\ion{H}{1} and CO velocities agree well, so do the H$\\alpha$ velocities.  In total, we cataloged 14 molecular clumps in NGC\\,4214. Our results from a virial mass analysis are compatible with a Galactic CO-to-H$_2$ conversion factor for NGC\\,4214 (lower than what is usually found in metal--poor dwarf galaxies). ",
        "introduction": "Previous studies have shown that molecular gas is hard to study in dwarf galaxies, even in those who are known to be rich in neutral atomic gas (\\ion{H}{1}), such as dwarf irregulars (dIrrs) and blue compact dwarfs (BCDs).  As carbon monoxide (CO) is the second most common molecular species after molecular hydrogen (H$_2$), and the one most easily observed in its cool, unshocked state, it is usually used as a tracer of molecular gas.  Relatively few dwarf galaxies have been detected in CO, though many have been observed (e.g., Taylor, Kobulnicky \\& Skillman 1998, Barone et al.\\ 2000).  Dwarf galaxies are set apart from from other \\ion{H}{1}-rich galaxies by their low metallicity: Recent work suggested that the CO-to-H$_2$ (`X') conversion factor is dependent upon the metallicity of the gas, with the conversion factor increasing as metallicity decreases (Arimoto, Sofue \\& Tsojimoto 1996, Wilson 1995, Verter \\& Hodge 1995, Rubio et al.\\ 1991).  Maloney \\& Black (1988) argued on theoretical grounds that at low metal abundances, the column density of CO will be lower, and there will be less self shielding from dissociating radiation.  This will cause the size of the CO emitting region to shrink, while leaving the H$_2$ unaffected.  Norman \\& Spaans (1997) predict that metallicities starting somewhere between 0.03 to 0.1 of the solar value are necessary before CO will be detectable.  Observationally, Taylor et al.\\ (1998) have found no dwarf galaxies with detectable CO emission at metallicities below $\\sim$ 0.15 solar. However, at present there are too few dwarfs with both metallicity measurements and CO observations to determine reliably if there is indeed a cutoff, and where exactly it might take place.  For the few dwarfs which have been detected in CO, following up those single dish detections with high spatial resolution interferometer data has proven enlightening.  In some cases, the high resolution of interferometers allow the direct estimate of the total mass of resolved giant molecular clouds in nearby dwarf galaxies independent of the CO intensity, and from this a first determination of the CO--to--H$_2$ conversion factor (e.g., Wilson 1995).  A list of dwarf galaxies which have been observed with interferometers to date is compiled in Tab.~1. E.g., Taylor et al.\\ (1999) used the Plateau de Bure interferometer to observe CO (1--0) and (2--1) in the post-starburst dwarf galaxy NGC~1569.  They found that the conversion factor was three times higher than in two other dwarf galaxies of the same metallicity (NGC\\,6822 and IC\\,10), suggesting that metallicity is not the only influence upon the conversion factor.  Motivated by this and other previous studies, we have observed another starburst dwarf galaxy, NGC~4214, at Owens Valley.  \\tablenum{1} \\placetable{table1} \\label{table1}  NGC~4214 is a nearby dwarf galaxy currently experiencing high levels of massive star formation.  Table~2 summarizes some of the properties of NGC~4214. Combining optical, NIR and UV data, Huchra et al.\\ (1983) concluded it went through a burst of star formation a few $\\times~10^7$ yr ago. NGC~4214 has two main regions of star formation -- a large complex of \\ion{H}{2} regions displaying a shell morphology in the galaxy center, and a smaller southern complex formed mainly of compact knots (MacKenty et al.\\ 2000). From age determinations of the \\ion{H}{2} regions and their spatial distribution, they argue that some of the younger star formation has been triggered by the interaction of the older \\ion{H}{2} regions with the surrounding ISM.  At a distance of 4.1 Mpc (Leitherer et al.\\ 1996), it is one of the closest starbursting dwarf galaxies, and thus an excellent target for a high resolution study of the molecular ISM.  Several groups have made single dish CO observations of NGC~4214 (Tacconi \\& Young 1985, Thronson et al.\\ 1988, Ohta et al.\\ 1993, Becker et al.\\ 1995, Taylor et al.\\ 1998). While the existence of CO emission away from the center of NGC~4214 can be inferred from the single dish data of Ohta et al.\\ (who, however, fail to detect the central emission peak) and Taylor et al., no analysis of the distribution of the molecular emission and its relation to the major star forming regions of the galaxy has been presented to date. Thronson et al.\\ (1988) have also measured the dust mass in a region about $4'$ in size using 160 $\\mu$m data.  The \\ion{H}{1} has been studied by Allsopp (1979) and McIntyre (1998) who find a large extended \\ion{H}{1} disk and a complex velocity field in the center. Based on the disturbed inner velocity field, Allsopp (1979) argued that NGC~4214 is interacting with NGC~4190, a dIrr with a projected distance of 35 kpc.  \\tablenum{2} \\placetable{table2} \\label{table2}  In Sec.~2 we describe our CO observations obtained with OVRO. In Sec.~3 we discuss the properties of the CO emitting regions detected by the interferometer, compare them to high--resolution VLA \\ion{H}{1} imaging as well as optical ground--based and HST imaging of NGC\\,4214. In Sec.~4 we will discuss the conversion factor, the star formation efficiency, implications on the star formation threshold as derived from our CO and \\ion{H}{1} data as well as the interplay of the star formation regions on the surrounding ISM in NGC\\,4214. We summarize our results in Sec.~5. ",
        "conclusions": "We have mapped the dwarf starburst galaxy NGC~4214 in the $^{12}$CO (1--0) transition at the OVRO millimeter array with a linear resolution of $\\sim 90$\\,pc. This first high-resolution study of the molecular component in this galaxy shows the presence of three molecular complexes which are in very different and distinct evolutionary stages. Both from the properties of the molecular gas and from HST and ground--based optical studies of the \\ion{H}{2}\\ regions and star clusters, the central region is the most evolved: Here, the molecular complex is highly disrupted, the impact of the massive star formation is largest, and a shell structure is in the process of forming. The star formation episode in the SE complex is more recent; thus, the dense ISM adjacent to the newly formed massive stars is far less disrupted, providing an ideal environment for ongoing, maybe self-propagating star formation proceeding further into the complex. Somewhat surprisingly, massive star formation has not started yet in the NW complex, possibly because the threshold volume density has not been reached in this region.  A comparison with the \\ion{H}{1} reveals that no clear correlation between the \\ion{H}{1} and CO surface densities exists. Although the NE and SW regions have almost the same total (\\ion{H}{1} and H$_2$) gas column density, violent star formation is only taking place in the SW region.  This nicely illustrates the irregular nature of these dwarf systems. The \\ion{H}{1} and H$\\alpha$ velocities are in good agreement with the CO velocities measured in the three molecular complexes. Assuming a Galactic conversion factor, the star formation efficiency for the SE complex is $\\sim 8$\\%, similar to the one found for spiral galaxies. This may be taken as an indication that the star formation process itself in dwarf galaxies proceeds in a similar way to that in spirals.  We were able to subdivide the molecular complexes further into 14 sub--complexes. Most of the clumps can not be resolved with our synthesized beam -- even the resolved objects are probably not GMCs. Deriving virial masses from these objects is not trivial and is affected by personal bias and the employed velocity and spatial resolution. Our clump analysis is compatible with a Galactic conversion factor for NGC\\,4214; this is substantially lower than the value found in other dwarf galaxies with similar metallicities so far. However, given the uncertainties in deriving this value we can not rule out the validity of the metallicity--X relation based on our data.  We derive a total molecular mass for NGC\\,4214 of 5.1$\\times10^6$\\,M$_{\\sun}$.  In summary, our OVRO observations of the molecular gas in the metal--poor dwarf galaxy NGC\\,4214 reveal an unexpected structural wealth -- they serve as a good excuse to map similar low--metallicity dwarf galaxies with interferometers in the future: even though the emission is faint when compared to spiral galaxies, high--resolution observations of the molecular gas can provide important insights to the irregular nature of dwarf galaxies."
    },
    "0011/astro-ph0011501_arXiv.txt": {
        "abstract": "This paper presents a series of global three dimensional accretion disk simulations carried out in the cylindrical limit in which the vertical component of the gravitational field is neglected.  The simulations use a cylindrical pseudo-Newtonian potential, $\\propto 1/(R-R_g)$, to model the main dynamical properties of the Schwarzschild metric.  The radial grid domain runs out to 60 $R_g$ to minimize the influence of the outer boundary on the inner disk evolution.  The disks are initially constant density with a Keplerian angular momentum distribution and contain a weak toroidal or vertical field which serves as the seed for the magnetorotational instability.  These simulations reaffirm many of the conclusions of previous local simulations.  The MRI grows rapidly and produces MHD turbulence with a significant Maxwell stress which drives accretion.  Tightly-wrapped low-$m$ spiral waves are prominent.  In some simulations radial variations in Maxwell stress concentrate gas into rings, creating substantial spatial inhomogeneities.  As in previous global simulations, there is a nonzero stress at the marginally stable orbit.  The stress is smaller than seen in stratified torus simulations, but nevertheless produces a small decline in specific angular momentum inside the last stable orbit.  Detailed comparisons between simulations are used to examine the effects of various choices in computational setup.  Because the driving instability is local, a reduction in the azimuthal computational domain to some fraction of $2\\pi$ does not create large qualitative differences.  Similarly, the choice of either an isothermal or adiabatic equation of state has little impact on the initial evolution.  Simulations that begin with vertical fields have greater field amplification and higher ratios of stress to magnetic pressure compared with those beginning with toroidal fields.  In contrast to MHD, hydrodynamics alone neither creates nor sustains turbulence. ",
        "introduction": "Recent increases in supercomputer performance have significantly improved the ability to evolve the basic equations of accretion disk structure and evolution.  These developments, along with continuing progress in understanding the most important physical processes that occur within accretion disks, suggest that predictive disk simulations are a realistic goal.  Such disk simulations will be global, fully three dimensional, and incorporate physical processes such as magnetohydrodynamics (MHD) and radiation transport.  At present, we are some ways from this goal; global simulations are still rather idealized in terms of disk structure, energetics, and dynamical range.  However, because almost any three dimensional disk simulation is relatively novel, there remain many significant questions to be investigated even with such simplified models.  To date there have been several global simulations of three dimensional disks, including Armitage (1998), Matsumoto (1999), Hawley (2000), Machida, Hayashi, \\& Matsumoto (2000), Hawley \\& Krolik (2001; hereafter HK), and Armitage, Reynolds \\& Chiang (2001; hereafter ARC). Much of this work has focused on thick accretion disks.  With a pressure scale height $H$ comparable to the disk radius $R$, the thick disk, or accretion torus, is more easily resolved in a numerical simulation than disks for which $H/R \\ll 1$.  Matsumoto (1999) followed the evolution of a thick torus embedded in an external vertical field, and found significant outflow collimated along the global vertical field lines.  Hawley (2000) considered tori containing toroidal fields and poloidal field loops, and Machida et al.~(2000) modeled a thick disk containing a toroidal field.  In these studies the initial field was entirely contained within the disk and the resulting outflows were confined to the creation of a magnetized corona.  A generic feature of all these thick disk simulations is the presence of large amplitude fluctuations in accretion rate, density, and other variables, in both space and time.  At a minimum these efforts have established that the magnetorotational instability, or MRI, (Balbus \\& Hawley 1991) is just as efficacious in thick disks as in local simulations to produce MHD turbulence and angular momentum transport.  Thick accretion tori with initially non-Keplerian angular momentum distributions are highly unstable.  MHD turbulence develops rapidly and is sustained by a self-consistent dynamo process within the disk.  The constant or near-constant specific angular momentum distribution of the initial torus rapidly evolves to one that is near Keplerian.  The main focus for dynamical studies would therefore seem to be Keplerian disks, both hot (high internal sound speed), and thin and cold (low internal sound speed).  Global simulations have been used to investigate specific physical issues in Keplerian accretion disk models.  HK examined the behavior of a thick, nearly Keplerian disk model to study the accretion flow through the radius of the marginally stable orbit ($r_{ms}$) in a pseudo-Newtonian potential.  They found significant stress at the location of the marginally stable orbit which creates a continuing decline in the value of the specific angular momentum $\\ell$ inside of this point.  The equation of state was adiabatic and there was no cooling, and the thickness of the disk remained roughly $H/R \\sim 0.15$ throughout.  So far there have been fewer simulations of Keplerian disks with low internal sound speeds corresponding to small $H/R$.  To sidestep the difficulty of resolving both $H$ and $R$, the cylindrical disk limit has been employed which omits vertical gravity and stratification. Armitage (1998) and Hawley (2000) simulated a few examples of these cylindrically-symmetric Keplerian disks.  More recently, ARC modeled several Keplerian cylindrical disks with sound speeds corresponding to $H/R$ of 0.08 and 0.04.  They also computed one stratified disk that covered a limited vertical extent.  As in the work of HK, they examined the inflow through $r_{ms}$.  Although they found a nonzero Maxwell stress inside of $r_{ms}$, it was not large enough to alter $\\ell$ to the same extent as seen in the simulation of HK.  They emphasized that the differences between their simulations and those of HK were quantitative; the same dynamical effects were present, only at reduced amplitude.  In addition to the lower initial sound speed and cylindrical symmetry, the ARC simulations used an angular domain of $\\pi/6$ compared with the full $2\\pi$ simulation of HK.  It is unclear which, if any, of these differences is the most significant in comparing the results of HK and ARC.  Clearly, we still have only a preliminary understanding of the specific processes that establish the precise turbulent stress levels in disks.  An attempt to investigate this question through simulation is further complicated by the need to distinguish between influences created purely by the numerical details (e.g.  computational domain, numerical resolution, initial conditions), and those determined by physics.  In an effort to characterize the array of technical choices in three dimensional simulations, this paper presents a series of disk simulations that begin with a Keplerian angular momentum distribution. These simulations are restricted to the cylindrical limit which omits the gravitational acceleration in the $z$ direction.  The cylindrical approximation has the distinct advantage of reducing the number of grid zones required in $z$:  only the wavelengths of the weak-field MRI need be resolved, not several pressure scale heights.  Here this advantage is exploited to increase the number of grid zones devoted to the radial extent of the disk, and its angular resolution.  Comparisons of cylindrical disks and vertically stratified disks in Hawley (2000) provided some evidence that the cylindrical disk is a good approximation for studying the evolution of density-averaged properties in a disk.  Here we will be exploring the detailed properties and limitations of these cylindrical disks in greater detail.  These cylindrical disks will be directly comparable to the simulations of ARC, and are the natural sequel to the local nonstratified shearing box simulations done previously (Hawley, Gammie, \\& Balbus 1995, hereafter HGB95; 1996; Matsumoto \\& Tajima 1995).  Comparison with shearing boxes provide insight into those phenomena that are truly global versus those that are well-described locally.  Cylindrical disk models will also provide a baseline of results that will be compared with future stratified Keplerian disk models.  Some of the questions that this investigation will address include: How do these global results compare with local shearing box simulations?  Are there effects that can be attributed to influences from the equation of state, or which arise when compared to simulations of tori with higher internal temperatures?  What differences are there in disks computed on grids with angular extent less that $2\\pi$?  Can a steady state be established in at least part of the disk?  What are the amplitudes, and the time- and length-scales for the fluctuations that result from the MRI-induced turbulence?  The plan of this paper is as follows.  In \\S 2 the computational setup is described.  Section 3 presents the results from a series of MHD cylindrical disk simulations.  In \\S 4 a purely hydrodynamic disk is considered for contrast.  The results and their consequences are discussed in \\S 5, and the conclusions are summarized in \\S 6. ",
        "conclusions": "Cylindrical disk simulations are a useful tool for investigating global evolution of disks evolving due to magnetically driven turbulence. Such simulations demonstrate that the conclusions developed in the local shearing box model hold in the global context as well.  As in the local model the MRI grows rapidly and produces MHD turbulence with a significant Maxwell stress.  The turbulence is more vigorous and more efficient in producing stress for a given total magnetic pressure when driven by an initial field that is vertical rather than toroidal. Hydrodynamics alone seems no more effective at creating or sustaining turbulence in a global model than it is in a local one.  In addition to reaffirming local properties of the MRI, cylindrical disk simulations illuminate disk characteristics that are truly global.  A net accretion rate is one such property, but there are also important nonlocal structural features.  Tightly-wrapped low-$m$ spiral waves are prominent.  The final accretion through the marginally stable orbit provides an example of a highly nonaxisymmetric spiral flow.  Particularly interesting in the present simulations is the tendency for radial variations in Maxwell stress to concentrate gas into rings, creating substantial spatial inhomogeneities.  A perennial question is the degree to which these simulations resemble traditional steady state $\\alpha$-disk models.  They do in so far as they accrete in direct response to internal stress, specifically due to MHD turbulence.  Beyond that, however, there are significant differences.  The simulations are characterized by large scale variability in space and time in all variables.  The stress is proportional to the magnetic pressure which is itself only indirectly related to other disk parameters.   In the simulations it is possible to approximate a quasi-steady state only with broad-stroke averages. In part this is due to the initial conditions (e.g., isolated tori or constant density slabs) which are far from a possible accreting steady state solution.  To address this issue it will be useful to attempt simulations that begin with more realistic initial states.  Results to date indicate, however, that analytic disk models are likely to prove woefully inadequate in describing detailed spatial and temporal disk properties so long as they are based upon a strict $\\alpha$ formulation with $\\alpha$ a constant in space and/or in time.  Although cylindrical disk simulations provide a valuable point of reference for future work, the lack of vertical stratification is clearly a major limitation for investigating many important physical processes.  This is obviously true for the development of a magnetized corona, or the launching of winds or jets.  It appears also to be an important factor in measuring the stress in the disk at the marginally stable orbit.  Stratified global thin disk simulations are the next logical step to contrast with existing global thick disk models.  Stratified thin simulations will require far more vertical grid zones centered around the equator than are used in cylindrical disks.  However, the tests presented here suggest that a reduction in the $\\phi$ domain is a acceptable problem simplification, as long as the potential for some small quantitative reduction in energy levels is kept in mind."
    },
    "0011/astro-ph0011392_arXiv.txt": {
        "abstract": "Thermal instability driven by radiation pressure might be relevant for intrinsically bright accreting sources.  The most promising candidate where this instability seems to be at work is one of the two known galactic superluminal sources, GRS 1915+105 (Belloni et al. 1997). In spite of being of relevance, this scenario has not yet been confirmed by proper time-dependent modelling. Non-linear time-dependent calculations performed by Szuszkiewicz and Miller (1998) show that thermally unstable discs undergo limit-cycle behaviour with successive evacuation and refilling of the central parts of the disc. This evolution is very similar to the one proposed by Belloni et al. (1997) in their phenomenological model. Further investigations are needed to confirm the thermal instability being operational in this source. First of all the spectra emitted from the disc during its evolution should be calculated and compared  with observations. Here such spectra are computed assuming local blackbody emission from the best studied transonic disc model. ",
        "introduction": "The simplest disc model which properly describes physical conditions around black holes is a transonic (slim) accretion disc. It is an axisymmetric, vertically integrated, non-self-gravitating, differentially rotating fluid configuration.  The assumption about self-gravity being negligible defines the outer edge of the disc and the requirement that the radial velocity equals the local sound speed somewhere in the disc, and then exceeds it, determines its radial structure. Once the mass, $M$, of the central black hole, accretion rate, $\\dot M$, and viscosity coefficient $\\alpha$ are specified the structure of the stationary disc is uniquely determined. However, such a disc might be subject to secular, thermal and pulsational instabilities. Here we focus on the thermal instability in which a temperature perturbation out of the equilibrium is enhanced by thermal processes. Moreover, we consider the thermal instability which is triggered  by one particular mechanism, namely  by the radiation pressure. The possibility that a global instability might develop in this model was indicated by the results of a local stability analysis. In the temperature - surface density diagram the sequence of equilibrium models forms a characteristic S-shape curve. Such curves reflect the  possibility of a transition between low and high accretion rate states, associated with an increase of the temperature. \\begin{figure}[h] % \\vspace{7.5cm} \\special{psfile=dots.eps hoffset=-30  voffset=-190 hscale=65 vscale=65 angle=0} \\caption[]{The phase portrait in the $\\log T$-$\\log\\Sigma$ plane for $r=10r_G$ (here $r_G$ is the Schwarzschild radius). The dotted line represents the equilibrium states, the so called S-curve. The lower branch of the S-curve overlaps with the trajectory of the evolving model. The continuous line shows the trajectory of the model during the evolution. The points marked 1 to 5 indicate those stages for which spectra are calculated. The points 1 to 4 correspond to the phase of disc evacuation, the point 5 is at the beginning of the refilling phase.} \\label{dot} \\end{figure} The first complete calculation of the global evolution of thermal instability in transonic discs was performed by Szuszkiewicz and Miller (1998). They have shown that at least for some values of $M$, $\\dot M$ and $\\alpha$ this evolution has a  cyclic character and consists of successive evacuation and refilling phases. An example of the limit-cycle for the best studied thermally unstable model with $M=10M_{\\odot}$, $\\dot M = 0.06$ (in Eddington units) and $\\alpha =0.1$ is shown in Figure 1. This particular non-stationary disc behaviour might have interesting observational consequences. ",
        "conclusions": "Transonic discs with accretion rates larger than a critical value which depends on the black hole mass and the viscosity parameter, show time dependent, cyclic behaviour due to the onset of thermal instabilities driven by radiation pressure. This behaviour can be described in terms of the cycle period, the duration of the bursting and quiescent phases, the amplitude and shape of the burst and finally by the spectral variability. For the different parameters of the system ($M$, $\\alpha$, $\\dot M$) we expect different variability properties and emitted spectra. Detailed theoretical predictions are definitely needed for modelling high-resolution timing and spectral data presently available for black hole systems."
    },
    "0011/astro-ph0011570.txt": {
        "abstract": "%%% YOUR CHANGES BELOW THIS LINE   A general overview of neutrino physics and astrophysics is given, starting with a historical account of the development of our understanding of neutrinos and how they helped to unravel the structure of the Standard Model. We discuss why it is so important to establish if neutrinos are massive and introduce the main scenarios to provide them a mass. The present bounds and the positive indications in favor of non-zero neutrino masses are discussed, including the recent results on atmospheric and solar neutrinos. The major role that neutrinos play in astrophysics and cosmology is illustrated. ",
        "introduction": "\\subsection{The hypothetical particle:}  One may trace back the appearance of neutrinos in physics to the discovery of radioactivity by Becquerel one century ago. When the energy of the electrons (beta rays) emitted in a radioactive decay was measured by Chadwick in 1914, it turned out to his surprise to be continuously distributed. This was not to be expected if the underlying process in the beta decay was the transmutation of an element $X$ into another one $X'$ with the emission of an electron, i.e. $X\\to X'+e\u0081\u00af$, since in that case the electron should be monochromatic. The situation was so puzzling that Bohr even suggested that the conservation of energy may not hold in the weak decays. Another serious problem with the `nuclear models' of the time was the belief that nuclei consisted of protons and electrons, the only known particles by then. To explain the mass and the charge of a nucleus it was then necessary that it had $A$ protons and $A-Z$ electrons in it. For instance, a $^4$He nucleus would have 4 protons and 2 electrons. Notice that this total of six fermions would make the $^4$He nucleus to be a boson, which is correct. However, the problem arouse when this theory was applied for instance to $^{14}$N, since consisting of 14 protons and 7 electrons would make it a fermion, but the measured angular momentum of the nitrogen nucleus was $I=1$.  The solution to these two puzzles was suggested by Pauli only in 1930, in a famous letter to the `Radioactive Ladies and Gentlemen' gathered in a meeting in Tubingen, where he wrote: `I have hit upon a desperate remedy to save the exchange theorem of statistics and the law of conservation of energy. Namely, the possibility that there could exist in nuclei electrically neutral particles, that I wish to call neutrons, which have spin 1/2 ...'. These had to be not heavier than electrons and interacting not more strongly than gamma rays.  With this new paradigm, the nitrogen nucleus became $^{14}$N$=14 p+7e+7`n$', which is a boson, and a beta decay now involved the emission of two particles $X\\to X'+e+`n$', and hence the electron spectrum was continuous. Notice that no particles were created in a weak decay, both the electron and Pauli's neutron $`n$' were already present in the nucleus of the element $X$, and they just came out in the decay. However, in 1932 Chadwick discovered the real `neutron', with a mass similar to that of the proton and being the missing building block of the nuclei, so that a nitrogen nucleus finally became just $^{14}$N$=7p+7n$, which also had  the correct bosonic statistics.  In order to account now for the beta spectrum of weak decays, Fermi called Pauli's hypotetised particle the neutrino (small neutron), $\\nu$, and furthermore suggested that the fundamental process underlying beta decay was $n\\to p+e+\\nu$. He wrote \\cite{fe34} the basic interaction by analogy with the interaction known at the time, the QED, i.e. as a vector$\\times$vector current interaction:  $$H_F=G_F\\int {\\rm d}^3x[\\bar\\Psi_p\\gamma_\\mu\\Psi_n] [\\bar\\Psi_e\\gamma^\\mu\\Psi_\\nu]+ h.c. .$$ This interaction accounted for the continuous beta spectrum, and from the measured shape at the endpoint Fermi concluded that $m_\\nu$ was consistent with zero and had to be small. The Fermi coupling $G_F$ was estimated from the observed lifetimes of radioactive elements, and armed with this Hamiltonian Bethe and Peierls \\cite{be34} decided to compute the cross section for the inverse beta process, i.e. for $\\bar\\nu+p\\to n+e^+$, which was the relevant reaction to attempt the direct detection of a neutrino. The result, $\\sigma=4(G_F^2/\\pi)p_eE_e\\simeq 2.3\\times 10^{-44}$cm$^2 (p_eE_e/m_e^2)$ was so tiny that they concluded `... This meant that one obviously would never be able to see a neutrino.'. For instance, if one computes the mean free path in water (with density $n\\simeq 10^{23}/$cm$^3$) of a neutrino with energy $E_\\nu=2.5$ MeV, typical of a weak decay, the result is $\\lambda\\equiv 1/n\\sigma\\simeq 2.5\\times 10^{20}$ cm, which is $10^7$AU, i.e. comparable to the thickness of the Galactic disk.  It was only in 1958 that Reines and Cowan  were able to prove that Bethe and Peierls had been too pessimistic, when they measured for the first time the interaction of a neutrino through the inverse beta process\\cite{re59}. Their strategy was essentially that, if one needs $10^{20}$ cm of water to stop a neutrino, having $10^{20}$ neutrinos a cm would be enough to stop one neutrino. Since after the second war powerful reactors started to become available, and taking into account that in every fission of an uranium nucleus the neutron rich fragments beta decay producing typically 6 $\\bar\\nu$ and liberating $\\sim 200$ MeV, it is easy to show that the (isotropic) neutrino flux at a reactor is $${d\\Phi_{\\nu}\\over d\\Omega}\\simeq {2\\times 10^{20}\\over 4\\pi} \\left( {{\\rm Power\\over GWatt}}\\right){\\bar\\nu\\over strad}.$$ Hence, placing a few hundred liters of water (with some Cadmium in it) near a reactor they were able to see the production of positrons (through the two 511 keV $\\gamma$ produced in their annihilation with electrons) and neutrons (through the delayed $\\gamma$ from the neutron capture in Cd), with a rate consistent with the expectations from the weak interactions of the neutrinos.   \\subsection{The vampire:}  Going back in time again to follow the evolution of the theory of weak interactions of neutrinos, in 1936 Gamow and Teller \\cite{ga36} noticed that the $V\\times V$ Hamiltonian of Fermi was probably too restrictive, and they suggested the generalization  $$H_{GT}=\\sum_iG_i[\\bar pO_in][\\bar eO^i\\nu]+h.c. ,$$ involving the operators $O_i=1,\\ \\gamma_\\mu,\\ \\gamma_\\mu\\gamma_5,\\ \\gamma_5,\\ \\sigma_{\\mu\\nu}$, corresponding to scalar ($S$), vector ($V$), axial vector $(A)$, pseudoscalar ($P$) and tensor ($T$) currents. However, since $A$ and $P$ only appeared here as $A\\times A$ or $P\\times P$, the interaction was parity conserving. The situation became unpleasant, since now there were five different coupling constants $G_i$ to fit with experiments, but however this step was required since some observed nuclear transitions which were forbidden for the Fermi interaction became now allowed with its generalization (GT transitions).  The story became  more involved when in 1956 Lee and Yang suggested that parity could be violated in weak interactions\\cite{le56}.  This could explain why the particles theta and tau had exactly the same mass and charge and only differed in that the first one was decaying to two pions while the second to three pions (e.g. to states with different parity). The explanation to the puzzle was that the $\\Theta$ and $\\tau$  were just the same particle, now known as the charged kaon, but the (weak) interaction leading to its decays violated parity.  Parity violation was confirmed the same year by Wu \\cite{wu57}, studying the direction of emission of the electrons emitted in the beta decay of polarised $^{60}$Co. The decay rate is proportional to $1+\\alpha \\vec{P}\\cdot \\hat{p}_e$. Since the Co polarization vector $\\vec P$ is an axial vector, while the unit vector along the electron momentum $\\hat{p}_e$ is a vector, their scalar product is a pseudoscalar and hence a non--vanishing coefficient $\\alpha$ would imply parity violation. The result was that electrons preferred to be emitted in the direction opposite to $\\vec{P}$, and the measured value $\\alpha\\simeq -0.7$ had then profound implications for the physics of weak interactions.  The generalization by  Lee and Yang of the Gamow Teller Hamiltonian was $$H_{LY}=\\sum_i[\\bar pO_in][\\bar eO^i(G_i+G_i'\\gamma_5)\\nu]+h.c. .$$ Now the presence of terms such as $V\\times A$ or $P\\times S$ allows for parity violation, but clearly the situation became even more unpleasant since there are now 10 couplings ($G_i$ and $G_i'$) to determine, so that some order was really called for.  Soon the bright people in the field realized that there could be a simple explanation of why parity was violated in weak interactions, the only one involving neutrinos, and this had just to do with the nature of the neutrinos. Lee and Yang, Landau and Salam \\cite{bright} realized that, if the neutrino was massless, there was no need to have both neutrino chirality states in the theory, and hence the handedness of the neutrino could be the origin for the parity violation. To see this, consider the chiral projections of a fermion $$\\Psi_{L,R}\\equiv {1\\mp \\gamma_5\\over 2}\\Psi.$$ We note that in the relativistic limit these two projections describe left and right handed helicity states (where the helicity, i.e. the spin projection in the direction of motion, is a constant of motion for a free particle), but in general an helicity eigenstate is a mixture of the two chiralities. For a massive particle, which has to move with a velocity smaller than the speed of light, it is always possible to make a boost to a system where the helicity is reversed, and hence the helicity is clearly not a Lorentz invariant while the chirality is (and hence has the desireable properties of a charge to which a gauge boson can be coupled). If we look now to the equation of motion for a Dirac particle as the one we are used to for the description of a charged massive particle such as an electron ($(i\\dslash-m)\\Psi=0$), in terms of the chiral projections this equation becomes $$i\\dslash\\Psi_L=m\\Psi_R$$ $$i\\dslash\\Psi_R=m\\Psi_L$$ and hence clearly a mass term will mix the two chiralities. However, from these equations we see that for $m=0$, as could be the case for the neutrinos, the two equations are decoupled, and one could write a consistent theory using only one of the two chiralities (which in this case would coincide with the helicity).  If the Lee Yang Hamiltonian were just to depend on a single neutrino chirality, one would have then $G_i=\\pm G_i'$ and parity violation would indeed be maximal. This situation has been described by saying that neutrinos are like vampires in Dracula's stories: if they were to look to themselves into a mirror they would be unable to see their reflected images.  The actual helicity of the neutrino was measured by Goldhaber et al. \\cite{go58}. The experiment consisted in observing the $K$-electron capture in $^{152}$Eu ($J=0$) which produced  $^{152}$Sm$^*$ ($J=1$) plus a neutrino. This excited nucleus then decayed into $^{152}$Sm ($J=0)+\\gamma$. Hence the measurement of the polarization of the photon gave the required information on the helicity of the neutrino emitted initially. The conclusion was that `...Our results seem compatible with ... 100\\% negative helicity for the neutrinos', i.e. that the neutrinos are left handed particles.  This paved the road for the $V-A$ theory of weak interactions advanced by Feynman and Gell Mann, and Marshak and Soudarshan \\cite{vma}, which stated that weak interactions only involved vector and axial vector currents, in the combination $V-A$ which only allows the coupling to left handed fields, i.e. $$J_\\mu=\\bar e_L\\gamma_\\mu\\nu_L+\\bar n_L\\gamma_\\mu p_L$$ with $H=(G_F/\\sqrt{2})J_\\mu^\\dagger J^\\mu$. This interaction also predicted the existence of purely leptonic weak charged currents, e.g. $\\nu+e\\to \\nu+e$, to be experimentally observed much later\\footnote{A curious fact was that the new theory predicted a cross section for the inverse beta decay a factor of two larger than the Bethe and Peierls original result, which had already been  confirmed in 1958 to the 5\\% accuracy by Reines and Cowan. However, in a new experiment in 1969, Reines and Cowan found a new value consistent with the new prediction, what shows that many times when the experiment agrees with the theory of the moment the errors tend to be underestimated.}.  The current involving  nucleons is actually not exactly $\\propto \\gamma_\\mu(1-\\gamma_5)$ (only the interaction at the quark level has this form), but is instead $\\propto \\gamma_\\mu(g_V-g_A\\gamma_5)$. The vector coupling remains however unrenormalised ($g_V=1$) due to the so called conserved vector current hypothesis (CVC), which states that the vector part of the weak hadronic charged currents ($J_\\mu^\\pm\\propto \\bar \\Psi\\gamma_\\mu \\tau^\\pm\\Psi$, with $\\tau^\\pm$ the raising and lowering operators in the isospin space $\\Psi^T=(p,n)$) together with the isovector part of the electromagnetic current (i.e. the term proportional to $\\tau_3$ in the decomposition $J^{em}_\\mu\\propto \\bar \\Psi\\gamma_\\mu (1+\\tau_3)\\Psi$) form an isospin triplet of conserved currents. On the other hand, the axial vector hadronic current is not protected from strong interaction renormalization effects and hence $g_A$ does not remain equal to unity. The measured value, using for instance the lifetime of the neutron, is $g_A=1.26$, so that at the nucleonic level the charged current weak interactions are actually ``$V-1.26A$''.  With the present understanding of weak interactions, we know that the clever idea  to explain parity violation as due to the non-existence of one of the neutrino chiralities (the right handed one) was completely wrong, although it lead to major advances in the theory and ultimately to the correct interaction. Today we understand that the parity violation is a property of the gauge boson (the $W$) responsible for the gauge interaction, which  couples only to the left handed fields, and not due to the absence of right handed fields. For instance, in the quark sector both left and right chiralities exists, but parity is violated because the right handed fields are singlets for the weak charged currents.  \\subsection{The trilogy:}  In 1947 the muon was discovered in cosmic rays by Anderson and Neddermeyer. This particle was just a heavier copy of the electron, and as was suggested by Pontecorvo, it also had weak interactions $\\mu +p\\to n+\\nu$ with the same universal strength $G_F$. Hincks, Pontecorvo and Steinberger showed that the muon was decaying to three particles, $\\mu\\to e\\nu\\nu$, and the question arose whether the two emitted neutrinos were similar or not. It was then shown by Feinberg \\cite{fe58} that, assuming the two particles were of the same kind, weak interactions couldn't be mediated by gauge bosons (an hypothesis suggested in 1938 by Klein). The reasoning was that if the two neutrinos were equal, it would be possible to join the two neutrino lines and attach a photon to the virtual charged gauge boson ($W$) or to the external legs, so as to generate a diagram for the radiative decay $\\mu\\to e\\gamma$. The resulting branching ratio would be larger than $10^{-5}$ and was hence already excluded at that time. This was probably the first use of `rare decays' to constrain the properties of new particles.  The correct explanation for the absence of the radiative decay was put forward by Lee and Yang, who suggested that the two neutrinos emitted in the muon decay had different flavour, i.e. $\\mu\\to e+\\nu_e+\\nu_\\mu$, and hence it was not possible to join the two neutrino lines to draw the  radiative decay diagram. This was confirmed at Brookhaven in the first accelerator neutrino experiment\\cite{da62}. They used an almost pure $\\bar\\nu_\\mu$ beam, something which can be obtained from charged pion decays, since the $V-A$ theory implies that $\\Gamma(\\pi\\to \\ell+\\bar\\nu_\\ell)\\propto m_\\ell^2$, i.e. this process requires a chirality flip in the final lepton line which strongly suppresses the decays $\\pi\\to e+\\bar\\nu_e$. Putting a detector in front of this beam they were able to observe the process $\\bar\\nu+p\\to n+\\mu^+$, but no production of positrons, what proved that the neutrinos produced in a weak decay in association with a muon were not the same as those produced in a beta decay (in association with an electron). Notice that although the neutrino fluxes are much smaller at accelerators than at reactors, their higher energies make their detection feasible due to the larger cross sections ($\\sigma\\propto E^2$ for $E\\ll m_p$, and $\\sigma\\propto E$ for $E\\gsim m_p$).  In 1975 the $\\tau$hird charged lepton was discovered by Perl at SLAC, and being just a heavier copy of the electron and the muon, it was concluded that a third neutrino flavour had also to exist. The direct detection of the $\\tau$ neutrino has recently been anounced by the DONUT experiment at Fermilab, looking at the short $\\tau$ tracks produced by the interaction of a $\\nu_\\tau$ emitted in the decay of a heavy meson (containing a $b$ quark) produced in a beam dump. Furthermore, we know today that the number of light weakly interacting neutrinos is precisely three (see below), so that the proliferation of neutrino species seems to be now under control.   \\subsection{The gauge theory:}  As was just mentioned, Klein had suggested that the short range charged current weak interaction could be due to the exchange of a heavy charged vector boson, the $W^\\pm$. This boson exchange would look at small momentum transfers ($Q^2\\ll M_W^2$) as the non renormalisable four fermion interactions discussed before. If the gauge interaction is described by the Lagrangian ${\\cal L}=-(g/\\sqrt{2})J_\\mu W^\\mu+h.c.$, from the low energy limit one can identify the Fermi coupling as $G_F=\\sqrt{2}g^2/(8M_W^2)$. In the sixties, Glashow, Salam and Weinberg showed that it was possible to write down  a unified description of electromagnetic and weak interactions with a gauge theory based in the group $SU(2)_L\\times U(1)_Y$ (weak isospin $\\times$ hypercharge, with the electric charge being $Q=T_3+Y$), which was spontaneously broken at the weak scale down to the electromagnetic $U(1)_{em}$. This (nowadays standard) model involves the three gauge bosons in the adjoint of $SU(2)$, $V_i^\\mu$ (with $i=1,2,3$), and the hypercharge gauge field $B^\\mu$, so that the starting Lagrangian is $${\\cal L}=-g\\sum_{i=1}^3J^i_\\mu V^\\mu_i-g'J^Y_\\mu B^\\mu + h.c. ,$$ with  $J^i_\\mu\\equiv \\sum_a \\bar \\Psi_{aL}\\gamma_\\mu (\\tau_i/2)\\Psi_{aL}$. The left handed leptonic and quark isospin doublets are $\\Psi^T=({\\nu_e}_L,e_L)$ and ($u_L, d_L)$ for the first generation (and similar ones for the other two heavier generations) and the right handed fields are SU(2) singlets. The hypercharge current is obtained by summing over both left and right handed fermion chiralities and is $J^Y_\\mu\\equiv \\sum_a Y_a\\bar \\Psi_{a}\\gamma_\\mu \\Psi_{a}$.  After the electroweak breaking one can identify the weak charged currents with $J^\\pm=J^1\\pm iJ^2$, which couple to the $W$ boson $W^\\pm=(V^1\\mp iV^2)/\\sqrt{2}$, and the two neutral vector bosons $V^3$ and $B$ will now combine through a rotation by the weak mixing angle $\\theta_W$ (with tg$\\theta_W=g'/g$), to give \\begin{eqnarray} \\pmatrix{ A_\\mu \\cr Z_\\mu}=\\pmatrix{ {\\rm c}\\theta_W & \\s\\theta_W\\cr -\\s\\theta_W & {\\rm c}\\theta_W} \\pmatrix{ B_\\mu\\cr V^3_\\mu}. \\end{eqnarray} We see that the broken theory has now, besides the massless photon field $A_\\mu$,  an additional neutral vector boson, the heavy $Z_\\mu$, whose mass turns out to be related to the $W$ boson mass through $s^2\\theta_W=1-(M_W^2/M_Z^2)$. The electromagnetic and neutral weak currents are given by $$J_\\mu^{em}=J^Y_\\mu+J^3_\\mu$$ $$J^0_\\mu=J^3_\\mu-\\s^2\\theta_W J^{em}_\\mu,$$ with the electromagnetic coupling being $e=g\\ \\s\\theta_W$.  The great success of this model came in 1973 with the experimental observation of the weak neutral currents using muon neutrino beams at CERN (Gargamelle) and Fermilab, using the elastic process $\\nu_\\mu e\\to \\nu_\\mu e$. The semileptonic processes $\\nu N\\to \\nu X$ were also studied and the comparison of neutral and charged current rates provided a measure of the weak mixing angle. From the theoretical side t'Hooft proved the renormalisability of the theory, so that the computation of radiative corrections became also meaningful.   \\begin{figure}[t] \\centerline{\\hbox{ \\psfig{figure=guanf1.eps} }} \\caption{\\footnotesize Neutral and charged current contributions to neutrino lepton scattering. \\label{eps1}} \\end{figure}   The Hamiltonian for the leptonic weak interactions $\\nu_\\ell+\\ell'\\to\\nu_\\ell+\\ell'$ can be obtained, using the Standard Model just presented, from the two diagrams in figure 1. In the low energy limit ($Q^2\\ll M_W^2,\\ M_Z^2$), it is just given by $$H_{\\nu_\\ell\\ell'}={G_F\\over \\sqrt{2}}[\\bar\\nu_\\ell \\gamma_\\mu (1-\\gamma_5)\\nu_\\ell][\\bar\\ell'\\gamma^\\mu(c_LP_L+c_RP_R)\\ell']$$ where the left and right couplings are $c_L=\\delta_{\\ell\\ell'}+ s^2\\theta_W-0.5$ and $c_R=s^2\\theta_W$. The $\\delta_{\\ell\\ell'}$ term in $c_L$ is due to the charged current diagram, which clearly only appears when $\\ell=\\ell'$. On the other hand, one sees that due to the $B$ component in the $Z$ boson, the weak neutral currents also  couple to the charged lepton right handed chiralities (i.e. $c_R\\neq 0$). This interaction leads to the cross section (for $E_\\nu\\gg m_{\\ell'}$) $$\\sigma(\\nu+\\ell\\to \\nu+\\ell)={2G_F^2\\over \\pi}m_\\ell E_\\nu\\left[c_L^2+{c_R^2 \\over 3}\\right],$$ and a similar expression with $c_L\\leftrightarrow c_R$ for antineutrinos. Hence, we have the following relations for the neutrino elastic scatterings off electrons $$\\sigma(\\nu_e e)\\simeq 9\\times 10^{-44}{\\rm cm}^2\\left( {E_\\nu\\over 10\\ {\\rm MeV}}\\right)\\simeq 2.5\\sigma(\\bar\\nu_e e) \\simeq 6\\sigma(\\nu_{\\mu,\\tau}e)\\simeq 7\\sigma(\\bar\\nu_{\\mu,\\tau}e).$$ Regarding the angular distribution of the electron momentum with respect to the incident neutrino direction, in the center of mass system of the process $d\\sigma(\\nu_e e)/d\\cos\\theta\\propto 1+0.1 [(1+\\cos \\theta)/2]^2$, and it is hence almost isotropic. However, due to the boost to the laboratory system, there will be a significant correlation between the neutrino and electron momenta for $E_\\nu\\gg $MeV, and this actually allows to do astronomy with neutrinos. For instance, water cherenkov detectors such as Superkamiokande detect solar neutrinos using this process, and have been able to reconstruct a picture of the Sun with neutrinos. It will turn also to be relevant for the study of neutrino oscillations that these kind of detectors are six times more sensitive to electron type neutrinos than to the other two neutrino flavours.  Considering now the neutrino nucleon interactions, one has at low energies (1~MeV$<E_\\nu<50$~MeV) $$\\sigma(\\nu_en\\to p e)\\simeq \\sigma(\\bar\\nu_ep\\to n e^+)\\simeq {G_F^2\\over \\pi}{\\rm c}^2\\theta_C(g_V^2+3g_A^2)E_\\nu^2,$$ where we have now introduced the Cabibbo mixing angle $\\theta_C$ which relates, if we ignore the third family, the quark flavour eigenstates $q^0$ to the mass eigenstates $q$, i.e. $d^0={\\rm c}\\theta_C d+\\s\\theta_Cs$ and $s^0=-\\s\\theta_C d+{\\rm c}\\theta_Cs$ (choosing a flavour basis so that  the up type quark flavour and mass eigenstates coincide).   \\begin{figure}[t] \\centerline{\\hbox{ \\psfig{figure=guanf2.eps} }} \\caption{\\footnotesize Neutrino nucleon and neutrino lepton cross sections (the three lines correspond, from top to bottom, to the $\\nu_e,\\ \\bar\\nu_e$ and $\\nu_{\\mu,\\tau}$ cross sections with electrons). \\label{sigma}} \\end{figure}  At $E_\\nu\\gsim 50$~MeV, the nucleon no longer looks like a point-like object for the neutrinos, and hence the vector ($v_\\mu$) and axial ($a_\\mu$) hadronic currents involve now momentum dependent form factors, i.e. $$\\langle N(p')|v_\\mu|N(p)\\rangle= \\bar u(p')[ \\gamma_\\mu F_V+{i\\over 2m_N}\\sigma_{\\mu\\nu}q^\\nu F_W] u(p)$$ $$\\langle N(p')|a_\\mu|N(p)\\rangle= \\bar u(p')[ \\gamma_\\mu \\gamma_5F_A+{\\gamma_5\\over 2m_N}q_\\mu F_P] u(p),$$ where $F_V(q^2)$ can be measured using electromagnetic processes and the CVC relation $F_V=F_V^{em,p}-F_V^{em,n}$ (i.e. as the difference between the proton and neutron electromagnetic vector form factors). Clearly $F_V(0)=1$ and $F_A(0)=1.26$, while $F_W$ is related to the magnetic moments of the nucleons. The $q^2$ dependence has the effect of significantly flattening the cross section. In the deep inelastic regime, $E_\\nu\\gsim$GeV, the neutrinos interact directly with the quark constituents. The cross section in this regime grows linearly with energy, and this provided an important test of the parton model. The main characteristics of the neutrino cross section just discussed are depicted in figure~2. For even larger energies, the gauge boson propagators enter into the play (e.g. $1/M_W^2\\to 1/q^2$) and the growth of the cross section is less pronounced  above 10~TeV ($\\sigma\\propto E^{0.36}$).  The most important test of the standard model came with the direct production of the $W^\\pm$ and $Z$ gauge bosons at CERN in 1984, and with the precision measurements achieved with the $Z$ factories LEP and SLC after 1989. These $e^+e^-$ colliders working at and around the $Z$ resonance ($s=M_Z^2 = ($91~GeV$)^2$) turned out to be also crucial for neutrino physics, since  studying the shape of the $e^+e^-\\to f\\bar f$ cross section near the resonance, which has the Breit--Wigner form $$\\sigma\\simeq{12\\pi \\Gamma_e\\Gamma_f\\over M_Z^2} {s\\over (s-M_Z^2)^2+ M_Z^2\\Gamma_Z^2},$$ it becomes possible to determine the total $Z$ width $\\Gamma_Z$. This width is just the sum of all possible partial widths, i.e. $$\\Gamma_Z=\\sum_f\\Gamma_{Z\\to f\\bar f}=\\Gamma_{vis}+\\Gamma_{inv}.$$ The visible (i.e. involving charged leptons and quarks) width $\\Gamma_{vis}$ can be measured directly, and hence one can infer a value for the invisible width $\\Gamma_{inv}$. Since in the standard model this last arises from the decays $Z\\to \\nu_i\\bar\\nu_i$, whose expected rate for decays into a given neutrino flavour is $\\Gamma_{Z\\to \\nu\\bar\\nu}^{th}=167$~MeV, one can finally obtain the number of neutrinos coupling to the $Z$ as $N_\\nu=\\Gamma_{inv}/ \\Gamma_{Z\\to \\nu\\bar\\nu}^{th}$. The present best value for this quantity is $N_\\nu=2.9835\\pm 0.0083$, giving then a strong support to the three generation standard model.  \\bigskip  Going through the history of the neutrinos we have seen that they have been extremely useful to understand the standard model. On the contrary, the standard model is of little help to understand the neutrinos. Since in the standard model there is no need for $\\nu_R$, neutrinos are  massless in this theory. There is however no deep principle behind this (unlike the masslessness of the photon which is protected by the electromagnetic gauge symmetry), and indeed in many extensions of the standard model neutrinos turn out to be massive. This makes the search for non-zero neutrino masses a very important issue, since it provides a window to look for physics beyond the standard model. There are many other important questions concerning the neutrinos which are  not addressed by the standard model, such as whether they are Dirac or Majorana particles, whether lepton number is conserved, if the neutrino flavours are mixed (like the quarks through the Cabibbo Kobayashi Maskawa matrix) and hence oscillate when they propagate, as many hints suggest today, whether they have magnetic moments, if they decay, if they violate CP, and so on. In conclusion, although in the standard model neutrinos are a little bit boring,  many of its extensions contemplate new possibilities which make the neutrino physics a very exciting field. ",
        "conclusions": ""
    },
    "0011/astro-ph0011447_arXiv.txt": {
        "abstract": "We study the spectra generated as a result of bulk Comptonization by relativistic electrons in radial flows onto compact objects. We solve numerically the general-relativistic radiative transfer equation in the Schwarzschild spacetime, in steady-state and under minimal assumptions. We show that power-law spectra result from multiple scatterings, in a way similar to thermal Comptonization. We also find that photon-electron interactions taking place near the black-hole event horizon affect very little the emerging spectra. We, therefore, argue that bulk Comptonization spectra do not carry distinguishing signatures of the compact object around which they are produced.  We examine the dependence of the spectra on simplifications often employed regarding the spacetime geometry, the distribution of photon sources, and the boundary conditions.  We show that the existence of trapped characteristics around a black hole reduces the efficiency of Comptonization and that general relativistic effects identically cancel bulk Comptonization effects for a free-falling flow and in the limit of infinitesimal mean-free path.  As a result, we find that neglecting the curved spacetime geometry leads to overestimating the high-energy flux by up to an order of magnitude.  Finally, we demonstrate that the spectrum from accretion onto a neutron star depends sensitively on the imposed boundary conditions, while that from a black hole is immune to such choices. ",
        "introduction": "\\label{sec:intro}  Compton scattering of photons in media with non-negligible bulk velocities is thought to be a significant, or even the dominant, radiative process in a number of astrophysical settings. Examples include the early Universe (the kinematic Sunyaev-Zeldovich effect; see, e.g., Rephaeli 1995), jets in blazars (see, e.g., Sikora, Begelman \\& Rees 1994), gamma-ray bursts (see, e.g., Lazzati et al.\\ 2000; Madau \\& Thompson 2000), and accretion flows onto galactic compact objects (see, e.g., Payne \\& Blandford 1981). The interaction of neutrinos with fast-moving electrons is also thought to determine the dynamics and fate of supernova explosions and super-critical accretion flows onto neutron stars (see, e.g., Fryer, Benz \\& Herant 1996; Burrows et al.\\ 2000).  Studying any of the above problems requires the solution of the kinetic equation that describes scattering of massless particles by fast-moving electrons, in which relativistic effects introduce very strong angular and photon-energy dependences. Because of such complexities, most previous studies have made a number of approximations in treating this radiative transfer problem (see Psaltis \\& Lamb 1997 and references therein). The kinetic equation is often expanded to different orders in the electron bulk velocity $\\beta$ (in units of the speed of light) and/or the photon energy $\\epsilon/m_{\\rm e}$ (where $m_{\\rm e}$ is the electron rest mass) and is truncated keeping only terms up to a low order \\cite{BP81,MK92,TMK97}. Furthermore, general relativistic effects are most usually neglected (but see Schmid--Burgk 1978; Zane et al.\\ 1996; Titarchuk \\& Zannias 1998; Laurent \\& Titarchuk 1999). Such analyses show that the effects of the odd- and even-order terms in $\\beta$ are qualitatively and quantitatively different and their relative importance is determined not only by the magnitude of the bulk velocity but by the properties of the radiation field and the flow as well (Psaltis \\& Lamb 1997, 2000).  The configuration in which Compton scattering of photons by relativistic electrons has been studied extensively in the past is that of accretion flows onto compact objects. Recent motivation for such studies comes from the observed power-law shapes and the absence of any detectable cut-offs (up to energies comparable to the rest mass of the electron) in the $\\gamma$-ray spectra of galactic accreting black holes (Grove et al.\\ 1998). Attributing these spectra to Comptonization by the relativistic electrons in the accretion flow could provide a spectral signature for the existence of a black hole in a galactic system and even lead to the measurement of its mass (as suggested in, e.g., Chakrabarti \\& Titarchuk 1996; Shrader \\& Titarchuk 1998, 1999; Laurent \\& Titarchuk 1999; Borozdin et al.\\ 1999).  Various analytical \\cite{PB81,MK92,TZZN96,TZ98,ZL00} and numerical (Zane et al.\\ 1996; Titarchuk et al.\\ 1997; Laurent \\& Titarchuk 1999; Psaltis 2000) treatments have been employed in previous studies of bulk Comptonization. However, the results of several analytic solutions have been recently questioned because of the inadequate order to which the transfer equations was truncated (see discussion in Psaltis \\& Lamb 1997, 2000) or of the inappropriate mathematical methods employed (see discussion in Zampieri \\& Lamb 2000).  Moreover, previous numerical studies have either considered a limited number of cases (as in, e.g., Zane et al.\\ 1996) or used approximate treatments which severely affected the properties of the solutions (see discussion in Psaltis \\& Lamb 2000).  With this paper we initiate a study of Compton scattering in relativistic accretion flows and winds, performed under a minimal set of assumptions and approximations, in order to address a number of questions that previous treatments raised. We solve numerically the kinetic equation for massless particles in a Schwarzschild spacetime derived by Lindquist (1966), using an iterative integration along the curved photon characteristics \\cite{S-B78,Zane96}. In this first study, we investigate the effects on the emerging spectra of properties related  to the flow velocity, the metric, the boundary conditions, and the distribution of sources in the flow. We only assume that the systematic down-scattering of photons as well as the Klein--Nishina corrections to the scattering cross section (which are of the same order) are negligible. Therefore, we cannot address issues related to the high-energy cut-offs of the spectra and their shapes at photon energies comparable to, or higher, than the electron rest mass; we will discuss these elsewhere. ",
        "conclusions": "We have solved the radiative transfer equation that describes scattering of photons by cold electrons moving at relativistic speeds, in steady-state, spherically symmetric accretion flows onto neutron stars and black holes.  We used an iterative procedure to integrate the radiative transfer equation along the photon characteristics. This is an efficient method that allows for the full description of the spacetime geometry and the velocity field (see Schmid--Burgk 1978 and Zane et al.\\ 1996), independent of the optical thickness of the flow. In this paper, we focused on examining simple configurations with the intention of exploring the generation of hard power-law spectra from accretion flows.  We have examined the dependence of the resulting spectra on the boundary conditions, the properties of the scattering medium, and the approximations often used for the spacetime geometry and the velocity field.  We demonstrated that the generation of power-law spectral tails is the result of multiple scatterings of the soft source photons by the fast moving cold electrons (see also Laurent \\& Titarchuk 1999), at both low and high optical depths.  As a result, the photons emerging from the accretion flow with high energies do not carry any signatures of regions of high electron velocities or spacetime curvature -contrary to expectations- but rather reflect the fact that a small number of photons can experience a significant number of scatterings even in a flow of small total optical depth.  This mechanism is very similar to the generation of power-law spectra by thermal Comptonization in static media. It differs, however, from the generation of power-law spectra by non-thermal electrons in optically thin media, which is the result of a single scattering by a power-law distribution of relativistic electrons (see Coppi 1999).  The multiple scatterings experienced by the photons tend to wipe out all memory of initial conditions such as the radial, angular, or energy distribution of the sources. However, the resulting spectra for {\\it neutron-star} flows do depend sensitively on the choice of the inner boundary conditions, confirming results reported earlier (Titarchuk et al.\\ 1997, Psaltis 1998; Psaltis \\& Lamb 2000). It is important to note, nevertheless, that none of the idealised boundary conditions used in our calculations is, strictly speaking, astrophysically relevant for a radial accretion flow onto a neutron star. Indeed, whether the photons that reach the inner boundary will be absorbed or reflected depends on the photon energy and the ionization state of the surface layers of the neutron star. Even if most of the photons are reflected, their energy gain at reflection will be determined by the thermal state and stratification of the transition region between the flow and the stellar surface. Moreover, the presence of shocks at this interface introduces additional complications. As a result, the spectral signature of a nearly-radial accretion flow onto a neutron star will not be uniquely determined by the flow properties but will be dictated by the specifics of the interaction of the flow with the stellar surface.  On the contrary, the spectral indices of the hard spectral tails resulting from the same flows onto black holes are immune to the inner boundary condition. This is due to the physical character of the inner boundary (the event horizon), i.e., that no radiation can escape to infinity (see Fig.~\\ref{fig:photon_escape} \\& \\ref{fig:ibc_effect}). For this reason, the spectra produced from bulk Comptonization in a {\\em purely radial, free-falling} flow onto a black hole have a power law index that is determined {\\it solely} by the optical depth (or mass accretion rate Eq.~[\\ref{eq:m_dot}]; see also discussion in Titarchuk \\ Zannias 1998). Furthermore, when compared to the spectral tails from flows onto neutron stars in this idealised limit, those of  black holes are steeper.  Several previous studies have examined the generation of power-law spectra without taking into account the curvature of the spacetime or the regime of highly relativistic speeds (e.g., Payne \\& Blandford 1981; Titarchuk et al.\\ 1987; Psaltis \\& Lamb 2000). We showed that, for a free-fall velocity profile, general relativistic effects identically cancel the bulk Comptonization effects {\\em everywhere\\/} in the flow, including radii very far from the compact object, as long as the photon mean-free path is very small.  If this were not true, then an observer comoving with the flow would be able to distinguish between being at rest at infinity and free-falling onto the compact object by making only local (because of the small photon mean-free path requirement) measurements of the evolution of the photon spectrum with time. However, for typical accretion rates, nowhere in the flow is the photon mean-free path small. Moreover, the photons that escape to infinity have experienced their last scatterings in regions of large photon mean-free path. As a result, the aforementioned cancellation of ${\\cal{O}}(\\beta^2)$ terms reduces the Comptonization efficiency --with respect to the flat geometry case-- only modestly in the case of flows onto neutron stars which are of relatively high optical depths.  Ignoring the actual spacetime geometry, when modelling a flow onto a black hole, results in more severe over-predicting of the hard tail flux, though.  For a black hole, the existence of trapped characteristics (see \\S\\ref{sec:numerics}) deprives the outer parts of the flow from a large fraction of the photons that have reached  close to the event horizon. Consequently, disregarding the Schwarzschild spacetime geometry can lead to overestimating the flux at high photon energies by more than an order of magnitude (e.g., at 100~$T_{\\rm in}$ in the case of Fig.~\\ref{fig:gr_flat_comparison}).  \\mbox{} We thank Fred Lamb, Feryal \\\"Ozel, Philipppos Papadopoulos, and Luca Zampieri for useful discussions and comments. This work was supported by a postdoctoral fellowship of the Italian  MURST (H.\\,P.), a postdoctoral fellowship of the Smithsonian Institution and also, in part, NASA (D.\\,P.).  We also thank MPI f\\\"ur Gravitationsphysik, Harvard-Smithsonian CfA, and University of Portsmouth (H.\\,P.) as well as SISSA (D.\\,P.) for their warm hospitality."
    },
    "0011/astro-ph0011136_arXiv.txt": {
        "abstract": "We present very deep and accurate photometry of the  open cluster M35. We have observed this association in the Cousins R,I filters, together with the Johnson  V filter. We have covered a region of   27.5$\\times$27.5 square arcmin, equivalent to a fifth of the total area of the cluster. The data range from I$_{\\rm c}$=12.5 to 23.5 magnitudes, and the  color intervals are 0.4$\\le$(V--I)$_{\\rm c}$$\\le$3.0, 0.5$\\le$(R--I)$_{\\rm c}$$\\le$2.5. Roughly, these values span  from  1.6 M$_\\odot$ down to the substellar limit, in the case of cluster members. By using the location of the stars on   color-magnitude and color-color  diagrams, we have selected candidate   members of this cluster. We have merged our sample with previously published data and obtained a color-magnitude diagram for the complete stellar population of the cluster, covering the spectral range early B -- mid M. Based on the distribution of field and cluster stars in color-magnitude and color-color diagrams, we estimate that two thirds of these candidates are likely to be true members of M35.  These stars approximately double the number of stars identified as candidate  members of this cluster ($\\sim$2700). We provide  the photometry and accurate positions of these stars. The deep photometry has allowed us to study the mass segregation within the cluster, the Luminosity Function  and Mass Function. We show that in the magnitude range 13 $\\le$ I$_c$ $\\le$ 22 there is a reduced  mass segregation, in opposition to what happens to higher mass stars, where the mass segregation is stronger. The Luminosity Function behaves essentially as the one characteristic of the Pleiades, presenting a peak at I$_{\\rm c}\\sim$19 magnitudes (M$_{\\rm I}\\sim$9). Combining our  photometry with previous data corresponding to more massive stars, we find that the Mass Function  increases monotonically, when plotted in a log-log form, until it reaches $\\sim$0.8 M$_\\odot$ ($\\alpha$=2.59). It remains shallower for  less massive stars ($\\alpha$=0.81 for 0.8--0.2 M$_\\odot$), whereas a decrease ins observed for stars close to the substellar regime. These different behaviors suggest that at least three mechanisms play a role in the formation of stellar and substellar objects. The total  mass of the cluster is $\\sim$1600 M$_\\odot$ in the area covered by this study. ",
        "introduction": "Open clusters serve as Rosetta stones to study stellar properties  because they provide homogeneous samples of known age, metallicity and distance. Therefore, open clusters can be used to determine empirically the evolution of different stellar properties such as the activity, rotation, chemical abundances, mass loss, et cetera. For low mass stars, however, only a handful of open clusters have been studied systematically, essentially due to the fact that most  clusters are far away, with distance moduli ranging from 6 to 12 magnitudes. With the recent improvements in detectors and with the availability of large ground-based telescopes, it is possible to obtain high quality data for a larger number of clusters.  M35 (NGC~2168) is a very interesting open cluster for several reasons. It is moderately nearby, with (m-M)$_0$=9.7 (Vidal 1973). It is very well populated, containing several thousand members with a total dynamical  mass estimated between 1600 and 3200 M$_\\odot$ (Leonard \\& Merritt 1989). It is spread over an area of $\\sim$1$\\times$1 square degrees, making it a  suitable target for multifiber spectroscopy. The contamination by field stars is not overwhelming, as happens in some other clusters of similar  characteristics. The reddening is also moderate --E(B--V)=0.17--0.225. Perhaps one of the most interesting properties is its youth, since it is said to be a cluster coeval with the Pleiades, based on  the fact that  their upper-main sequence turn-off locations are said to be  at the same position in color-magnitude diagrams (Vidal 1973).  Until very recently, only the upper-main sequence  (MS) of M35 had been studied with some detail. Photometric (UBV filters) and/or proper motion searches were carried out by Cuffey (1938), Hoag et al. (1961),  Cudworth (1971), Vidal (1973), McNamara \\& Sekiguchi (1986a). They reached barely V=15, about  1.2 M$_\\odot$, in the case of M35 members. At the present time, different groups have conducted several  multi-band  CCD photometry which reach  much deeper members, such as Sung \\& Bessell (1999=SB1999), von Hippel et al. (2000) and Sarrazine et al. (2000). The last works present data in the  UBVRI filters, reaching beyond V=19, whereas Sung \\& Bessell (1999) used the  UBVI passbands only, reaching one magnitude fainter.  Sung \\& Bessell (1999) were able to estimate the interstellar reddening, the distance and the  age of the cluster, deriving E(B-V)=0.255$\\pm$0.024, (m-M)$_0$=9.60$\\pm$0.10 and $\\tau$=200$_{-100}^{+200}$ Myr. Previously, Sung \\& Lee (1992) derived a distance modulus (m-M)$_0$=9.3, an age $\\tau$=85 Myr, and a differential internal extinction E(B-V)=0.26--0.44. A different estimate of the distance and reddening have been published by von Hippel et al. (2000), E(B-V)=0.3, (m-M)$_V$=10.25. These authors also estimated an age by fitting isochrones, yielding $\\tau$=100 Myr. However, they inferred that M35 is older than the Pleiades by comparing in a differential way members of both clusters. Based on an extension of the data presented by von Hippel et al. (2000), and assuming a metallicity of [Fe/H]=--0.21 (Barrado y Navascu\\'es et al. 2000a,b), Sarrazine et al. (2000)  derived (m-M)=10.16$\\pm$0.10, E(B--V)=0.198$\\pm$0.008 and $\\tau$=160$\\pm$40 Myr. This age and the Sung \\& Bessell (1999) value are significantly  older than the canonical upper-main sequence age of the Pleiades, 70-100 Myr (Mermilliod 1981; Meynet, Mermilliod, \\& Maeder 1993). These older ages are in agreement with our own findings concerning the distribution of rotational velocities and lithium abundances (Barrado y Navascu\\'es et al. 2000a, 2000b). Based on all the available information, they assumed an age of $\\tau$=175 Myr. Note that Barrado y Navascu\\'es et al. (2000b) have derived the first metallicity estimate for M35 based on high resolution spectroscopy, with the resulting value of [Fe/H] = --0.21$\\pm$0.10, lower than the value corresponding to the Pleiades, [Fe/H]=+0.01, estimated in the same fashion by these studies. The Pleiades cluster is one of the cornerstones of astrophysics, and the comparison of  members of both clusters can provide additional information for understanding their properties. Although several   stellar characteristics are thought to depend essentially on age (rotation, coronal and chromospheric activity, et cetera), this assumption has been recently challenged by Randich \\& Schmitt (1995), comparing the X-ray luminosities of Praesepe  and the Hyades, two coeval clusters (see also the discussion in Barrado y Navascu\\'es et al. 1997, 1998a). Extensive and detailed studies of several clusters of similar age  can help to understand this crucial problem, as well to clarify the evolution of the stellar properties when comparing clusters of different ages.    We have embarked on an ambitious study  of the open cluster M35, combining high resolution spectroscopy to determine lithium abundances, radial and rotational velocities (Barrado y Navascu\\'es et al. 2000a, 2000b) with deep imaging to extend the membership list to nearly the hydrogen burning mass limit (or beyond) and estimate the cluster Mass Function (Barrado y Navascu\\'es et al. 1998b). The current paper is devoted to the presentation of our deep imaging data, to study the complete MF,  the mass segregation within the cluster and to estimate the total mass of M35. ",
        "conclusions": "We have obtained deep V(RI)$_{\\rm c}$ photometry of the open cluster M35, reaching much fainter stars than the previous surveys. The comparison of the location of the stars in color--magnitude and color-color diagrams allowed us to establish the preliminary membership of the detected stars. We have also estimated the contamination by field stars. We have analyzed this accurate, deep  photometry of M35. Our survey has shown that in the studied color range there is a weak mass  segregation, in opposition to the stronger segregation present for more massive stars. These data have allowed us to establish the Luminosity Funtion and Mass Function of the low and very low mass components of the clusters. The LF presents a peak at I$_{\\rm c}$=19 magnitudes, corresponding to  a mass of $\\sim$0.3 M$_\\odot$. This result is very similar to the Pleiades' LF. By merging our own dataset with data from  previous studies of the cluster, which correspond to more massive stars, we have been able to obtain a complete mass function for the stellar component of the cluster, covering the mass range 6.0 $\\le$ Mass $\\le$ 0.08 M$_\\odot$. When expressed in a  log-log form, the MF increases monotonically in the 6-0.8 M$_\\odot$ domain with a power index of $\\alpha$=$+$2.59. After this point, the power index is reduced significantly, to a value of $\\alpha$=$+$0.81. Below this point, the behavior changes and the power law index become negative ($\\alpha$=--0.8). Finally, we have estimated the total mass of the cluster in the area were our data were collected. After taking into account the pollution by field stars, the total mass contained in the central part of the cluster is $\\sim$1600 M$_\\odot$. It would be very interesting to collect deeper photometry of the cluster, reaching well below the substellar limit, to verify whether the M35 MF behaves as the Pleiades and sigma Orionis Cluster or if a significant fraction of the M35 BDs have already been dissipated from the cluster."
    },
    "0011/astro-ph0011070_arXiv.txt": {
        "abstract": "Hubble's Law, $v=HD$ (recession velocity is proportional to distance), is a theoretical result derived from the Friedmann-Robertson-Walker metric. $v=HD$ applies at least as far as the particle horizon and in principle for all distances. Thus, galaxies with distances greater than $D=c/H$ are receding from us with velocities greater than the speed of light and superluminal recession is a fundamental part of the general relativistic description of the expanding universe. This apparent contradiction of special relativity (SR) is often  mistakenly remedied by converting redshift to velocity using SR. Here we show that galaxies with recession velocities faster than the speed of light are observable and that in all viable cosmological models, galaxies above a redshift of three are receding superluminally. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011246_arXiv.txt": {
        "abstract": "We report on extremely high velocity molecular gas, up to $-80$ km s$^{-1}$ relative to the ambient medium, in the giant star-formation complex 30 Doradus in the Large Magellanic Cloud (LMC), as observed in new 22 GHz H$_2$O $6_{16}\\rightarrow5_{23}$ maser emission spectra obtained with the Mopra radio telescope. The masers may trace the velocities of protostars, and the observed morphology and kinematics indicate that current star formation occurs near the interfaces of colliding stellar-wind blown bubbles. The large space velocities of the protostars and associated gas could result in efficient mixing of the LMC. A similar mechanism in the Milky Way could seed the galactic halo with relatively young stars and gas. ",
        "introduction": "The most luminous H {\\sc ii} region in the Local Group of galaxies is 30 Doradus (Tarantula nebula, N157: Kennicutt 1984; Walborn 1991) in the Large Magellanic Cloud (LMC). It is the site of recent star formation on a massive scale and harbours a rich stellar cluster (NGC2070) of which the core (R136: Feast et al.\\ 1960) has only recently been resolved into many individual stars (cf.\\ Hunter 1999). The H {\\sc ii} region is embedded in an extended halo of expanding gas. In between, the stellar winds from the rapidly evolving massive stars interact with each other and with the surrounding interstellar medium (ISM) to create an intricate complexity of shells, filaments and dense knots (Meaburn 1980, 1984; Wang \\& Helfand 1991; Chu \\& Kennicutt 1994). The radiation from these hot stars photo-evaporates the surrounding dusty, molecular and neutral clouds (Hunter et al.\\ 1995; Poglitsch et al.\\ 1995; Scowen et al.\\ 1998). It thus appears that the newly born stars destroy their parent Giant Molecular Cloud (GMC), which might prevent further star formation. Yet the recent discovery of protostars and structures similar to Herbig-Haro objects (Jones et al.\\ 1986; Walborn \\& Blades 1987, 1997; Hyland et al.\\ 1992; Rubio et al.\\ 1992, 1998; Walborn et al.\\ 1999) provides ample evidence for the occurrence of current star formation, in spite of --- or perhaps due to (Dopita 1981; Caulet et al.\\ 1982; Dopita et al.\\ 1985) --- the violent interaction between colliding stellar-wind blown bubbles.  The first evidence of ongoing star formation in 30 Doradus was the detection of H$_2$O maser emission from 0539$-$691 (Whiteoak et al.\\ 1983), a region $\\sim1.5^\\prime$ NE of R136. H$_2$O masers have been found in three more regions in the LMC (Scalise \\& Braz 1982; Whiteoak et al.\\ 1983; Whiteoak \\& Gardner 1986) amongst which is 0540$-$696 in the GMC extending $\\sim3$ kpc due S of 30 Doradus (Israel 1984; Cohen et al.\\ 1988; Kutner et al.\\ 1997; Johansson et al.\\ 1998). H$_2$O masers occur in dust-enshrouded (Class 0: Shu et al.\\ 1987) Young Stellar Objects (YSOs) and Ultra Compact H {\\sc ii} regions (e.g.\\ Knapp \\& Morris 1976; Genzel \\& Downes 1977; Cesarsky et al.\\ 1978; Rodriguez et al.\\ 1978; Codella et al.\\ 1994, 1996), and disappear once the H {\\sc ii} region becomes diffusive after $\\sim10^5$ yr (Genzel \\& Downes 1979; Codella \\& Felli 1995). H$_2$O maser emission arises both in low-mass and massive YSOs (Haschick et al.\\ 1983; Wouterloot \\& Walmsley 1986; Tereby et al.\\ 1992; Wilking et al.\\ 1994; Codella et al.\\ 1995) but is more intense in massive cores of star formation (Palla et al.\\ 1991; Henning et al.\\ 1992; Palagi et al.\\ 1993; MacLeod et al.\\ 1998). Models (Elitzur et al.\\ 1989, 1992) suggest that the H$_2$O masers occur in shocked layers at the neutral side of photo-dissociation regions (PDRs) around newly formed stars. Masers are extremely bright emission lines, that can be used both to probe the kinematics within the circumstellar environment as well as to trace the systemic velocity of the masing object.  Here we present new 22 GHz data of the known H$_2$O maser source 0539$-$691 and a hitherto unknown H$_2$O maser at $\\sim2.5^\\prime$ SE of R136. Both show discrete emission components at large blue-shifted velocities with respect to the bulk of the ISM in the 30 Doradus region. This places the masers and hence the star formation sites at the rim of expanding stellar-wind blown bubbles, providing new evidence for sequential, shock-induced star-formation. ",
        "conclusions": "The 22 GHz detections of 0539$-$691 \\& 0539$-$691B show clear morphological and, for the first time, kinematical evidence for the location of H$_2$O masers where rapidly expanding volumes of ionized gas collide either with other such gas volumes or with a neutral and/or molecular medium. The stellar-wind blown bubbles have supersonically swept up material from the surrounding ISM, which then became highly compressed and shocked. Upon the collision of such bubble with another bubble or with a dense cloud, this dense fast-moving layer becomes even more shocked and compressed. This may be a favourable condition for the formation of new stars (Dopita 1981; Caulet et al.\\ 1982).  On kinematic grounds we can at least distinguish three masers: two in 0539$-$691 and one in 0539$-$691B. The large velocity difference between the 182 km s$^{-1}$ maser and the other maser emission in 0539$-$691 strongly suggests also a large spatial separation and the location in distinct structures in the ISM. The exact location of the H$_2$O masers inside of the dense molecular layer may be either near the PDR around the protostar and/or near the PDR associated with the expanding shell. Until interferometric observations at 22 GHz are undertaken it is not yet possible to identify the masers with a particular structure as seen at other wavelengths. The reason is that the morphological and kinematical structure of the ISM within $\\sim100$ pc from R136 is extremely complex on all scales.  Protostar candidates in the 30 Doradus complex have been associated with dark globules that are being photo-evaporated externally (Clayton 1987; Rubio et al.\\ 1992; Scowen et al.\\ 1998; Walborn et al.\\ 1999) and are situated in between expanding gas shells (Hyland et al.\\ 1992). These shells range in size from $\\sim0.5$ pc around individual massive stars (Hunter et al.\\ 1995), up to the $\\sim50$ pc size shells around associations of $\\sim10^2$ OB and/or Wolf-Rayet stars (Chu \\& Kennicutt 1994). In fact, both the size distribution of structures seen in H {\\sc i} (Stanimirovi\\'{c} et al.\\ 1999) and the size-turbulence relation of H {\\sc ii} regions (Roy et al.\\ 1986) obey a Kolmogorov scaling law. The co-existence of ionized and molecular gas in the PDRs at the interface of the shells and molecular clouds indicates highly localized (clumpy) and stratified physical structures (Poglitsch et al.\\ 1995; Scowen et al.\\ 1998). Poglitsch et al.\\ estimate typical clump sizes of a few pc, containing a mass of $M\\sim10^{2-3}$ M$_\\odot$. This is just below the detection limit of the currently available CO surveys. Assuming a star formation efficiency of $\\sim10$\\% and a Salpeter-law Initial Mass Function (Salpeter 1955), the most massive star formed in a $10^3$ M$_\\odot$ cloud is expected to have $M\\sim8$ to 10 M$_\\odot$ --- i.e.\\ an early-B type star with a luminosity of $L\\sim10^4$ L$_\\odot$. Hence one or a few bright H$_2$O masers may be expected in such a star forming cloud.  It seems thus the case that star formation as traced by H$_2$O maser emission occurs at several locations in the 30 Doradus complex, which are associated with fast-moving gas ramming into other material. This confirms the picture of the propagation of sequential star formation, in which stellar-wind blown bubbles are now starting off a new epoch of star formation in 30 Doradus. It remains to be proven whether, in a similar fashion, the formation of the young stellar population of $\\sim10^6$ yr --- that includes R136 and that is responsible for the H {\\sc ii} region of 30 Doradus as we see it today --- was induced by a previous star formation event $\\sim2\\times10^7$ yr ago as traced by red supergiants within 45 pc from R136 (Hyland et al.\\ 1992).  The star formation propagating throughout the 30 Doradus complex may initially have been triggered by ram pressure resulting from the supersonic motion of the LMC through the extended galactic halo (de Boer et al.\\ 1998): the LMC appears to have been moving from apo-galacton at $\\sim100$ to 200 kpc one Gyr ago to peri-galacton at $\\sim50$ kpc today (Gardiner et al.\\ 1994; Heller \\& Rohlfs 1994). A close encounter between the Magellanic Clouds $\\sim2$ to $5\\times10^8$ yr ago (Gardiner et al.\\ 1994; Heller \\& Rohlfs 1994) may also have played a r\\^{o}le in stirring up the ISM in the LMC, triggering enhanced star formation as evidenced by common peaks in the age distributions of clusters in both Magellanic Clouds (Pietrzy\\'{n}ski \\& Udalski 2000). More may be learnt about the triggering of sequential star formation from studies of the 3 kpc long reservoir of CO gas due S of 30 Doradus (Israel 1984; Cohen et al.\\ 1988; Kutner et al.\\ 1997) that seems to be in an earlier stage of star formation activity than the 30 Doradus complex (Heydari-Malayeri et al.\\ 1999).  If the newborn stars have space velocities similar to those of the H$_2$O masers, then these stars will quickly mix with the older stellar populations of the LMC. Together with the ejection of stars by supernovae in binaries (Zwicky 1957; Blaauw 1961; Kaper et al.\\ 1997; Hoogewerf et al.\\ 2000) and dynamical interactions within young clusters (Poveda et al.\\ 1967; Gies \\& Bolton 1986; Hoogewerf et al.\\ 2000), high-velocity star formation may thus provide an additional mechanism for the formation of runaway stars. In fact, it has been argued that the star cluster associated with SN 1987A has formed at a velocity of 50 km s$^{-1}$ relative to the disk (Graff \\& Gould 2000). Such high velocity corresponds to 5 kpc in $10^8$ yr, sufficient to cross the bar of the LMC. Thus both gas and young stars may efficiently mix throughout the entire galaxy.  There is some indication that this may also happen in the Milky Way: the presence of early-type Main-Sequence stars in the galactic halo (Rolleston et al.\\ 1999) may be explained by their formation in gas entities that were being ejected out of the galactic disk. Some of this gas may fall back towards the galactic disk in the form of high-velocity clouds (Richter et al.\\ 1999). We see the galactic disk edge-on, but the LMC disk nearly face-on, which makes it much easier to detect this escaping gas and associated star formation in the LMC than in our own Milky Way galaxy."
    },
    "0011/astro-ph0011520_arXiv.txt": {
        "abstract": "We report the detection of $^{12}$CO~($1 \\rightarrow 0$) and $^{12}$CO~($2 \\rightarrow 1$) emission from the central regions ($\\lesssim 5$--$10 {\\rm \\ kpc}$) of the two powerful radio galaxies 3C~31 and 3C~264.  Their individual CO emission exhibits a double-horned line profile that is characteristic of an inclined rotating disk with a central depression at the rising part of its rotation curve.  The inferred disk or ring distributions of the molecular gas is consistent with the observed presence of dust disks or rings detected optically in the cores of both galaxies.  For a CO to H$_2$ conversion factor similar to that of our Galaxy, the corresponding total mass in molecular hydrogen gas is $(1.3 \\pm 0.2) \\times 10^9 {\\rm \\ M_{\\odot}}$ in 3C~31 and $(0.31 \\pm 0.06) \\times 10^9 {\\rm \\ M_{\\odot}}$ in 3C~264.  Despite their relatively large molecular-gas masses and other peculiarities, both 3C~31 and 3C~264, as well as many other powerful radio galaxies in the (revised) 3C catalog, are known to lie within the fundamental plane of normal elliptical galaxies.  We reason that if their gas originates from the mergers of two gas-rich disk galaxies, as has been invoked to explain the molecular gas in other radio galaxies, then both 3C~31 and 3C~264 must have merged a long time (a few billion years or more) ago but their remnant elliptical galaxies only recently (last tens of millions of years or less) become active in radio.  Instead, we argue that the cannibalism of gas-rich galaxies provides a simpler explanation for the origin of molecular gas in the elliptical hosts of radio galaxies.  Given the transient nature of their observed disturbances, these galaxies probably become active in radio soon after the accretion event when sufficient molecular gas agglomerates in their nuclei. ",
        "introduction": "Bright radio sources served as the first signposts for highly energetic activity in galaxies.  The nature of these galaxies, and the reason for their luminous radio activity, have since been subjects of detailed investigation.  Although the vast majority resemble luminous elliptical galaxies, observations by  \\citet{hec86} and \\citet{smi89a,smi89b} showed that a significant fraction of the most powerful radio galaxies at low redshifts (henceforth meaning $z \\lesssim 0.1$) exhibit peculiar optical morphologies suggestive of close encounters or mergers between galaxies.  Moreover, \\citet*{smi90} showed that the visibly disturbed galaxies exhibit a slightly greater degree of rotational support than normal elliptical galaxies, and hence also are dynamically peculiar.  Nevertheless, all of the powerful radio galaxies examined by \\citet{smi90}, mostly selected from the 3C catalog, lie within the fundamental plane of normal elliptical galaxies.  Some radio galaxies possess so much dust that they can be detected in the far-infrared by {\\it IRAS} \\citep*{gol88}.  This discovery has motivated many subsequent searches for molecular gas in radio galaxies, in all cases by selecting those with known appreciable amounts of dust.  At low redshifts, eleven radio galaxies have so far been detected in $^{12}$CO \\citep*{ric77,phi87,laz89,mir89,maz93,eva99,eva99a,eva99b}.  The detected galaxies span nearly three orders of magnitude in radio luminosity, and comprise those exhibiting core-dominated radio sources as well as classical double-lobed FR-I (edge darkened) and FR-II (edge brightened) radio sources.  All have inferred molecular-gas masses between $10^{9} {\\rm \\ M_{\\odot}}$ and $10^{11} {\\rm \\ M_{\\odot}}$, except for the very nearby radio galaxy Centaurus~A which has a molecular-gas mass of $\\sim$$2 \\times 10^{8} {\\rm \\ M_{\\odot}}$ \\citep{eck90}.  In four of the five cases (neglecting 3C~71, which is a spiral galaxy) mapped, Cen~A \\citep{phi87,eck90}, B2~1318+343 \\citep*{rad91}, 3C~293 \\citep{eva99a}, and 4C~12.50 \\citep{eva99b}, the CO gas is found to be concentrated in a compact (diameter of a few kpc or smaller) ring or disk around the center of the galaxy; in the fourth case, 3C~84 (Perseus~A), the gas is asymmetrically distributed over an extended (over 10~kpc) region, but also coinciding with the center of the galaxy \\citep{ino96}.  The molecular gas in radio galaxies may therefore comprise the reservoir for fueling their central supermassive black holes.  At low redshifts, the only other type of galaxy to commonly exhibit molecular gas masses as high as $\\ga 10^{10} {\\rm \\ M_{\\odot}}$ are infrared-luminous galaxies.  These gas and dust rich galaxies exhibit vigorous star formation thought to be triggered by galaxy-galaxy interactions, and indeed the majority of ultraluminous infrared galaxies (i.e., those with $L(60{\\mu}{\\rm m}) \\ga 10^{12} L_{\\odot}$) are found to be merging systems of gas-rich disk galaxies.  In the latter systems, the CO gas is found to be preferentially concentrated in a ring or disk around the nuclear regions of the merging galaxies \\citep{dow98,bry99}.  The observed similarities have prompted the suggestion that radio galaxies also originate from the mergers of two gas-rich disk galaxies, and in many cases comprise their still disturbed E/SO merger products \\citep{mir89,maz93,eva99a,eva99b}.  Given the predisposition of the abovementioned surveys towards relatively dust-rich objects, are the radio galaxies so far detected biased towards those with unusually large amounts of molecular gas?  All the initial detections (except for 3C~71 and Centaurus~A) were made with the NRAO 12-m telescope, and in many cases reached lower limits just below the minimum gas mass detected.  As a result, it is not clear whether the IR-bright radio galaxies so far detected represent the most gas-rich members of a population that all possess substantial amounts of dust and gas, or extreme members of a population that possess a broad range of dust and gas masses.  This issue is of importance for a proper understanding of the nature of radio galaxies, and what fuels their supermassive black holes.  To address this issue, we have initiated a deep survey of all the previously undetected radio galaxies at redshifts $z \\le 0.031$ in the revised 3C catalog \\citep{spi85}.  The objects in this catalog represent the most luminous radio galaxies --- the majority of which exhibit classical double-lobed radio morphologies --- at their particular redshifts in the northern hemisphere; the selected objects therefore represent the most luminous radio galaxies in the Local Universe.  At low redshifts the vast majority can be clearly seen to be luminous elliptical galaxies, which in most cases comprise first ranked galaxies in poor clusters of roughly ten members although a number reside in much richer environments \\citep[][and references therein]{zir96,zir97}.  Only a small minority have been detected in the infrared by {\\it IRAS}; indeed, only three objects in this catalog have previously been detected in CO, one of which lies in the redshift range of our survey. ",
        "conclusions": ""
    },
    "0011/astro-ph0011185_arXiv.txt": {
        "abstract": "We construct a simple model for stationary, axisymmetric black-hole magnetospheres, in which the poloidal magnetic field is generated by a toroidal electric current in a thin disk with the inner edge, by solving the vacuum Maxwell equations in Schwarzschild background. In this work, to obtain a concise analytical form of the magnetic stream function, we use the approximation that the inner edge is far distant from the event horizon. The global magnetospheric structure with the closed-loop and open field lines threading the inner and outer parts of the disk is explicitly shown, claiming that the model is useful as a starting point to study astrophysical problems involving inward disk-driven winds to a black hole and outward ones to infinity. The asymptotic shape of the field lines at the event horizon becomes nearly cylindrical, while at infinity it becomes conical. The magnetic spot in the disk connected with the black hole through the loop field lines occupies a very narrow region with the ring area roughly equal to the horizon area. By taking account of the existence of a uniform (external) magnetic field, we also obtain the model for collimated open field lines. Then, it is found that the magnetic connection between the black hole and the disk breaks down if the uniform field is strong enough. Considering slow rotation of the magnetosphere and angular momentum transfer by inward winds from the disk, the final discussion is devoted to gradual disruption of the closed loops due to radial accretion of disk plasma toward the black hole. ",
        "introduction": "It is widely believed that a supermassive black hole surrounded by an accretion disk can work as a central engine of active galactic nuclei. If magnetic field exists around the black hole and the disk, it is important to understand the magnetospheric structure for explaining the highly energetic phenomena. The magnetohydrodynamical (MHD) approach in general relativity may be necessary to study various features of the black hole magnetospheres. Unfortunately, even if the magnetosphere is assumed to be stationary and axisymmetric, any global models describing both inward winds in the strong-gravity region near black hole and outward winds in the far distant region have not been constructed, in spite of a large number of works devoted to the MHD problems. This is mainly because it is very difficult to analyze the highly nonlinear Grad-Shafranov equation for the magnetic stream function which defines the poloidal field lines (see e.g., \\citet{bes97} for review).  To avoid the mathematical difficulty in the full MHD system, one may consider vacuum solutions of the magnetic stream function as a practical approximation applicable to the magnetically dominated models. In particular, the well-known examples such as split-monopole, uniform and paraboloidal fields in Schwarzschild geometry \\citep{mi73,w74,bz77} have been used as the zeroth-order approximation for the perturbation method in slow-rotation and force-free fields, from which the first-order perturbation has been derived for giving a weak toroidal magnetic field and discussing the Blandford-Znajek process \\citep{bz77}. Another important MHD problem to be studied is fluid motion along poloidal field lines. For examples, the existence of accretion flows onto the black hole passing through some MHD critical points has been shown under a fixed shape of field lines \\citep{tak90}. The vacuum solutions allowing a field configuration of astrophysical interest will become a useful tool in an attempt to discuss the trans-critical motion of outgoing and ingoing magnetized winds, except in the region where the effect of fluid inertia should dominate to change crucially the field structure.  Of course, even for the vacuum fields, some current distribution should be assumed to exist in a restricted region of the magnetosphere. In this paper we focus on the magnetic field generated by a toroidal current distribution in a thin disk around the Schwarzschild black hole. The source of the above-mentioned fields with split-monopole and paraboloidal structures can be also disk currents. However, such models contain only open field lines and fail to describe a magnetic connection of the disk with the black hole. This will be partially due to the absence of the disk inner edge separated from the event horizon. The magnetic connection by closed field lines is expected to exist and have astrophysically important effects \\citep{pu91a,gru99,bl00,li00c,kl00}. Hence, our purpose here is to present explicitly a global magnetospheric model with both closed and open field lines threading a thin disk and connecting magnetically the inner and outer parts, respectively, with the black hole and the far distant region, which is relevant to the problems of inward and outward disk-driven winds. Though the vacuum model would miss some important features of more realistic magnetospheres, our modeling could be a preliminary step for understanding the various MHD processes in further detail.  To obtain general models based on a disk current distribution, it may be convenient to consider linear superposition of the magnetic field generated by a single ring current located at a fixed radius on the equatorial plane which has been found in Kerr geometry \\citep{pet75}. Assuming the current distribution from an inner edge to an outer one, \\citet{li00a} has derived a formal expression of the magnetic stream function in terms of the infinite sum of the multipole fields, without revealing the global shape of the field lines. Because the each multipole field of which the structure has been discussed by \\citet{gos00} still fails to give closed field lines regular at the event horizon, the infinite sum is essential to understanding the global shape.  In \\S2 we would like to develop a different approach to arrive at the infinite sum formula for the magnetic stream function. In \\S3 we succeed in reducing it to an analytical simple form by locating the disk inner edge at a radius much larger than the horizon radius and the outer edge at infinite radius. The global structure of closed and open field lines is shown, and it is found that only the very narrow region of the disk separated from the inner edge can be magnetically connected with the black hole, together with the result that the field lines become nearly cylindrical at the event horizon and conical at infinity. In \\S4 we also consider a modification of the global shape due to an addition of a uniform external field. Though collimated field lines appear at the far distant region, it is found that for a strong uniform field the magnetic connection between the black hole and the disk breaks down. Finally we discuss the MHD effect due to slow rotation of the magnetosphere (namely, angular momentum transfer by magnetically dominated inward winds from the disk to the black hole), which could lead to radial accretion of disk plasma and gradual disruption of the closed field lines \\citep{pc89}. In the following we use the geometrical units with $G=c=1$. ",
        "conclusions": "If the plasma inertia is taken into account even under the condition of magnetic domination, the asymptotic collimation from the conical shape of field lines should occur in logarithmic scales of the cylindrical radius \\citep{chi91,tom94}. Apparently such an inertial back-reaction of the plasma on the poloidal field does not work in the vacuum fields. Then, we present a model with collimated structure, by adding the uniform field (which is a solution of the vacuum Maxwell equations) \\begin{equation} \\Psi_{u} \\ = \\ \\pi B_{u}r^{2}\\sin^{2}\\theta \\end{equation} to $\\Psi$ given by equation (41). Let us denote the new magnetic stream function by $\\Psi'(=\\Psi+\\Psi_{u})$. Because $\\Psi$ is negative in the far distant region, $B_{u}$ should be chosen to be negative for allowing a smooth collimation of the conical field lines. If the uniform field is weak, namely $B_{0}+B_{u}>0$, the value of $\\Psi'$ remains positive at the event horizon to keep the magnetic connection between the disk and the black hole (Fig.~3). However, if $B_{0}+B_{u}$ becomes negative, the field lines threading the black hole can reach to infinity only (Fig.~4).    Of course, this addition of a uniform field is just a mathematical procedure, and the strong collimation shown in Figures~3 and~4 is not a consequence of the physical confinement mechanisms by self-generated toroidal fields and/or by external pressure with a boundary at finite cylindrical radius. It remains quite uncertain whether the details of the field-line shape shown in these figures become approximately valid in magnetically dominated wind regions or not. Nevertheless, it is interesting to note that the vacuum models can successfully describe the transition of the structure of field lines threading the black hole (i.e., from magnetic connection with the disk to that with the remote load), which may occur in more realistic magnetospheres as an observable change of the astrophysical activity.  To discuss the magnetic activity by using the vacuum field $\\Psi'$ (or $\\Psi$) as a background field perturbed by magnetically dominated winds from the disk, let us consider the magnetosphere slowly rotating with the angular velocity $\\Omega_{F}$ of a magnetic field line and the hole's one $\\Omega_{H}$ (i.e., $r\\sin\\theta\\Omega_{F}\\ll 1$ and $M\\Omega_{H}\\simeq a/4M\\ll 1$). If a magnetic field line threading the black hole is closed in the disk at $r=r_{d}$ (see Figs.~1 and~3), the angular velocity $\\Omega_{F}$ is nearly equal to the local Keplerian one, i.e., \\begin{equation} \\Omega_{F} \\ \\simeq \\ \\Omega_{K} \\ \\equiv \\ \\left( M/r_{d}^{3}\\right)^{1/2} \\ , \\end{equation} and we have $r_{d}\\simeq M\\chi w_{0}$, where $\\chi$ is a numerical factor in the range $1<\\chi\\leq 1.9$ for the $s=2$ model. Now the approximation of slow rotation is assured by virtue of the estimation such that $r_{d}\\Omega_{F}\\simeq1/\\sqrt{\\beta w_{0}}\\ll1$. However, the inner edge in realistic accretion disks cannot be too distant from the event horizon (namely, $w_{0}$ cannot be too large), and from the relation \\begin{equation} \\Omega_{H}/\\Omega_{F} \\ \\simeq \\ M\\Omega_{H}(\\beta w_{0})^{3/2} \\ , \\end{equation} we expect to be $\\Omega_{H}<\\Omega_{F}$ as a typical case for slowly rotating black holes. Hence, the inward winds from the magnetic spot carry angular momentum from the disk to the black hole, and the disk plasma which remains in the magnetic spot without flowing into the magnetic tube will be torqued and fall inward. This means that the disk currents located in the inner region and sustaining the loop field lines become unstable for radial accretion, and the closed loop field disconnected from the disk must disappear owing to the infinite redshift effect at the event horizon \\citep{pc89}.  In the magnetically-dominated wind region the toroidal magnetic field $B_{T}$ is approximately given by \\begin{equation} B_{T} \\ \\simeq -4\\pi\\eta L/\\alpha\\varpi \\ , \\end{equation} where $\\alpha$ and $\\varpi$ are the lapse function and the cylindrical radius, respectively, in Kerr geometry. By virtue of the approximation of slow rotation we obtain $\\alpha=\\sqrt{1-(2M/r)}$ and $\\varpi=r\\sin\\theta$. Along a poloidal field line the particle flux $\\eta$ per unit flux tube and the total angular momentum $L$ of magnetized fluid are conserved, and the angular momentum flux per unit area carried by the winds towards the black hole is given by $-\\eta LB_{p}$, which is a consequence of exerting the magnetic torque on the disk plasma. (Now $L$ is positive, while $\\eta$ is negative for ingoing winds.) Then, the characteristic timescale $t_{1}$ to extract the plasma angular momentum in the disk would be \\begin{equation} t_{1} \\ \\sim \\ h\\rho_{d}r_{d}^{2}\\Omega_{K}/(-\\eta LB_{p}) \\ \\sim \\ (r_{d}/h)/\\beta_{e}\\Omega_{K} \\ , \\end{equation} where $\\rho_{d}$ is the mass density in the disk with the thickness $h$, and $\\beta_e$ is the ratio of the external magnetic stress $B_{T}B_{p}$ at the disk surface to the local pressure in the disk. One may also consider the internal torque transporting angular momentum radially outward through the disk, for which the accretion timescale $t_{2}$ is estimated to be \\begin{equation} t_{2} \\ \\sim \\ (r_{d}/h)^{2}/\\beta_{i}\\Omega_{K} \\ , \\end{equation} and the ratio $\\beta_{i}$ of the internal stress to the local pressure is usually supposed to be less than unity (e.g., $\\beta_{i}\\simeq 0.1$). In our modeling the external torque responsible for carrying angular momentum into the winds should be dominant if compared with the internal torque, and we require the condition $\\beta_{e}\\gg\\beta_{i}h/r_{d}$, which becomes valid for $h\\ll r_{d}$, unless the turbulent internal disk field is supposed to be much stronger than any ordered vertical field at the disk surface. Hence, for $\\beta_{e}\\approx\\beta_{i}$, the duration of the closed loop structure could be quite longer than the dynamical timescale $\\Omega_{K}^{-1}$ particularly in a thin disk.  In significant long term accretion of plasma beyond the timescale $t_{1}$ one must consider the change of the closed field-line topology, unless dynamo action in the inner disk region can efficiently work to recover the closed loops. The process of disruption of the closed loops would be quasi-stationary, and our models could give a qualitative picture of the structure change as follows: Initially strong disk currents can sustain nearly circular loops of field lines shown in Figure~1 (i.e., $B_{0}\\gg|B_{u}|$). As the disk currents become weaker (i.e., as the ratio $B_{0}/|B_{u}|$ decreases) by virtue of the radial accretion, the loops begin to collapse towards the black hole (especially in vertical direction), as was shown in Figure~3. Finally, only the loops disconnected from the horizon remains near the inner edge (see Fig.~4 for $B_{0}\\leq|B_{u}|$). The loop field completely annihilates as a set of O-points, if the radial accretion from the inner edge continues still more \\citep{pc90b}. During this process, a part of electromagnetic energy stored in loop field lines could be outward transferred as a Poynting flux to infinity. We expect that the magnetic field (41) to become a convenient initial state for analytical and numerical MHD calculations to study the problem of non-steady energy release in plasma accretion toward the black hole (see, e.g., \\citet{ko00} for numerical investigations, in which the initial magnetic field is chosen to be the Wald solution in Kerr background).  In the slowly rotating magnetospheres the closed field line plays the role of a path for the energy and angular momentum transfer from the disk to the black hole. [ \\citet{li00c} claimed that the transfer from the black hole to the disk occurs only for $a/M>0.36$.]  However, the spin-up of the black hole should stop when the structure shown in Figure~4 has been established. Then, the energy extraction from the rotating black hole becomes possible as a mechanism of energy release. The angular velocity $\\Omega_{F}$ of the open field lines connecting the black hole with the remote load may satisfy the condition $\\Omega_{F}\\approx\\Omega_{H}/2$, which allows the twisted field with the negative toroidal component $B_{T}$ (i.e., $\\eta<0$ and $L<0$) to bring negative angular momentum and energy to the black hole. The efficiency of this Blandford-Znajek mechanism is still controversial, because the validity of the condition $\\Omega_{F}\\approx\\Omega_{H}/2$ is based on the assumption that a causal contact between the black hole and the remote plasma holds. In plasma accretion onto the black hole the poloidal inward velocity exceeds the fast magnetosonic speed near the horizon, and no information could propagate outward across some critical surface \\citep{pc90a}. In fact, the black hole-driven winds satisfying $\\Omega_{F}\\ll\\Omega_{H}$ have been found to exist \\citep{pu98}. For such winds the main mechanism of energy extraction from the black hole could be the screw instability, which makes the twisted magnetic tubes unwind and causes a sudden transport of the magnetic free-energy into the plasma's kinetic energy. For the cylindrical magnetic field shown in Figure~4 we can use the instability condition known as the Kruskal-Shafranov criterion \\citep{kad66} of the form \\begin{equation} |B_{T}|/2\\pi\\delta \\ > \\ |B_{p}|/S \\ , \\end{equation} where $\\delta$ and $S$ are the width and the length of the cylindrical magnetic tube. This has been proved by \\citet{gru99} in the force-free magnetosphere, but it should be noted that the analysis has been limited to the case such that $\\Omega_{F}=0$ and $S\\gg M$ (i.e., the approximation of no rotation and no gravity). Then, a naive application \\citep{li00b} of the Kruskal-Shafranov criterion to the magnetosphere where $B_{T}$ is roughly estimated to be $r\\sin\\theta\\Omega_{F}B_{p}$ seems to be very dubious. Here, the toroidal field is estimated by equation (46), which remains valid even for $\\Omega_{F}=0$, and considering the region well away from the event horizon (i.e., $\\alpha\\simeq1$), we rewrite the instability condition into the form \\begin{equation} S \\ > \\ |B_{p}|\\delta^{2}/(2|\\eta L|) \\ , \\end{equation} of which the right-hand side is constant if the magnetic flux $|B_{p}|\\delta^{2}$ in the cylinder is conserved. This means that if a larger angular momentum flux $|\\eta LB_{p}|\\delta^{2}$ is carried away from the black hole by magnetically dominated winds in the cylinder with the width $\\delta\\sim M$, the twisted field lines become screw-unstable in a shorter scale $S$. After the disruption of the screw-unstable structure of magnetic field, the rotating black hole would begin to twist again the magnetic field lines, and the energy release from the black hole could quasi-periodically occur.  As was first proposed by \\citet{gru99}, the closed field lines connecting the black hole with the disk may be also screw-unstable. However, for the field lines closed in the disk the toroidal component will be mainly generated by disk rotation with the angular velocity $\\Omega_{K}\\simeq\\Omega_{F}$. Then, we must discuss the instability condition by including the effect of $\\Omega_{F}$, which is a problem studied in future works \\citep{ma01}.  Finally it should be emphasized that the closed field-line topology presented in this paper is a very transient structure in rapidly rotating plasma-filled magnetospheres. The angular velocity $\\Omega_{F}\\simeq\\Omega_{K}$ of closed loops connecting the black hole with the disk will be smaller than the hole's angular velocity $\\Omega_{H}$, and the energy and angular momentum transfer from the black hole to the disk could spin up the disk plasma \\citep{li00c}. Then, as a consequence of this energy extraction from the black hole, outward flows of plasma from the disk would be generated, and owing to the frozen-in condition the field lines threading the black hole are eventually disconnected from the disk. The closed field lines anchored in plasma orbiting in the ergosphere would also open up to infinity by the strong dynamo action, as was conjectured to occur in field-aligned pulsars \\citep{pc90b}. Further, the timescale $t_{2}$ for radial accretion due to internal torque may become shorter than the timescales to open up the loops closed in the disk. Then, the closed loops accrete toward the black hole and annihilate owing to the infinite redshift effect \\citep{pc90b}. These are astrophysically interesting processes of rapid energy release, and we could observe transient flaring states of black hole magnetospheres, if the repeated formation of closed magnetic loops due to strong disk currents in the inner region occurs."
    },
    "0011/astro-ph0011472_arXiv.txt": {
        "abstract": "We study star formation history and stellar metallicity distribution in galaxies in a $\\Lam$ cold dark matter universe using a hydrodynamic cosmological simulation. Our model predicts star formation rate declining in time exponentially from an early epoch to the present with the time-scale of 6 Gyr, which is consistent with the empirical Madau plot with modest dust obscuration. Star formation in $L^*$ galaxies continues intermittently to the present also with an exponentially declining rate of a similar time-scale, whereas in small galaxies star formation ceases at an early epoch. The mean age of the extant stars decreases only slowly with increasing redshift, and exceeds $1\\Gyr$ at $z=3$. Normal galaxies contain stars with a wide range of metallicity and age: stars formed at $z<1$ have metallicity of $0.1-1.0\\Zsun$, while old stars take a wide range of values from $10^{-6}\\Zsun$ to $3.0\\Zsun$. The mean metallicity of normal galaxies is in the range $0.1-1.0\\Zsun$. Dwarf galaxies that contain only old stars have a wide range of mean metallicity ($10^{-4}-1.0\\Zsun$), but on average they are metal deficient compared with normal galaxies. ",
        "introduction": "\\label{intro_section} Over the last 15 years, cold dark matter (CDM) models have served as basic frameworks to study the formation of cosmic structure, and have been successful in delineating a scenario of galaxy formation \\citep{Davis85, Blumenthal84}. We now understand the formation of large-scale structure reasonably well in terms of a CDM model dominated by a cosmological constant $\\Lam$ \\citep{Efstathiou90, Ostriker95, Turner97, Balbi00, Lange00, Hu00}.   When and how galaxies formed, however, are a far more complicated problem due to non-gravitational physical processes operating on small scales. Many authors use semianalytic models to study galaxy formation, where dark matter halo formation under the hierarchical structure formation is supplemented heuristically with physical processes for baryons \\citep{White91, Baugh98, Kauffmann99a, SP99, Cole00}. They have succeeded to give a picture of galaxy formation roughly consistent with observations, but have the disadvantage that  one has to assume a set of simplified model equations, often associated with additional free parameters, for each physical process included in the model. The advantage of semianalytic models is, on the other hand, that they are computationally light, and therefore, one can search for a viable model with a small amount of cost.   An alternative but computationally much more expensive approach is direct cosmological hydrodynamic simulations. So far, intergalactic medium and overall galaxy formation processes have been well studied with hydrodynamic approach \\citep[\\eg,][]{CO92a, CO93, % CO00, Katz96} \\citep[see][however]{Pearce99}. Much effort invested in improving the accuracy of simulations has brought the hydrodynamic mesh approaching to $(1000)^3$, % with which we may hope that we obtain meaningful results on some aspects of galaxy properties. It is certainly impossible to resolve internal structure of galaxies with the present simulations. Nevertheless, we would expect that global properties such as global star formation rate (SFR) are described reasonably well with a few adjustable parameters, since they primarily depend on thermal balance of the bulk of clouds, such as how gas cools, or how surrounding gas is reheated by feedback by star formation. The fact that we obtain a reasonable amount of baryonic mass frozen into stars and a global metal abundance with reasonable input parameters justifies, at least in part, our expectation.   In this paper, we primarily discuss the star formation history of galaxies and the evolution of stellar metallicity, which are the direct output of the simulation. The aspects that require additional use of a stellar population synthesis model, such as luminosity function and colors of galaxies, will be discussed in a separate publication \\citet[][hereafter Paper II]{Nagamine01}. Since the hydrodynamic approach is computationally demanding, we do not attempt to make a fine tuning of input parameters, but our aim here instead is to make qualitative predictions which can be used to either verify or falsify the $\\Lam$CDM model by future observations with only a limited number of assumptions, treating the dynamics of baryons more accurately than in semianalytic models.   Our results would give insight on how galaxies are assembled from an early epoch to the present. Observationally interesting questions would include: (i) How old are stars in normal and dwarf galaxies, and what is their age distribution? A relevant question is whether star formation takes place continuously or intermittently;  (ii) What is the stellar metallicity distribution in normal and dwarf galaxies? A relevant old problem is the paucity of metal poor stars in the solar neighborhood \\citep[G-dwarf problem:][]{Bergh62, Schmidt63}; (iii) Is there a unique age-metallicity relation for stars in normal galaxies, and (iv) Is there a relation between the mass of galaxies and average metallicity? How these aspects evolve with redshift is also an observationally relevant problem, although the observational studies could answer only for global properties for high-redshift objects. Our simulation yield unambiguous predictions to these problems, at least qualitatively.    In \\S~\\ref{simulation_section}, we describe our simulation. In \\S~\\ref{sf_section} we first present the global star formation history of the entire simulation box. We then discuss in \\S~\\ref{divide_sf_section} the star formation history of individual galaxies. The mean age of the stars and the formation epoch of galaxies are discussed in \\S~\\ref{age_section}. We discuss the stellar metallicity distribution and the mean metallicity of galaxies in \\S~\\ref{metal_section}, and conclude in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We have studied predictions of the $\\Lam$CDM universe with regard to the star formation history and the stellar metallicity distribution in galaxies. Our purpose has been to clarify what observations would test the validity of the model of galaxy formation based on $\\Lambda$CDM. We have first shown that the global SFR averaged over all galaxies declines with a characteristic time-scale of $\\tau\\approx 6\\Gyr$. About a quarter of stars formed by $z\\approx 3.5$ and another quarter by $z\\approx 2$. This star formation history is consistent with the empirical Madau plot if modest dust obscuration is assumed.   Our calculation indicates that star formation in $L^*$ galaxies continues {\\it intermittently} to the present epoch as they accrete gas and merge with smaller systems with a decline rate corresponding to an e-folding time of $\\tau\\approx 6\\Gyr$. Star formation in less massive galaxies ceases earlier with an e-folding time of $\\tau \\ltsim 1\\Gyr$. In particular, dwarf galaxies cease their star formation by $z=2$. It is not excluded that minor star formation activity at low-redshift is missed in our calculation due to yet insufficient resolution, but it is unlikely that the global trend would be modified. We consider that this is a feature generic to the CDM structure formation scenario, and observational tests would provide a vital test.   Stars formed in the CDM galaxy formation scenario do not follow a unique age-metallicity relation, but show a considerable spread of metallicity for a given age of stars. Metallicity of very young stars spreads from $0.1-1.0\\Zsun$. The spread gradually increases towards high redshift. It takes $0.01-1.0\\Zsun$ at $z=2$. At $z=3$, the distribution spreads over from $10^{-6}\\Zsun$ to super-metal rich ($3\\Zsun$). Young stars are necessarily metal rich, but old stars are not necessarily metal poor. The variation of the mean metallicity with cosmic time is only gradual. The average metallicity for galaxies at $z\\approx 1$ differs little from that of local galaxies.   On the other hand, mean metallicity varies as a function of the mass of galaxies. Less massive galaxies are metal poor. Another interesting prediction is that mean metallicity of dwarf galaxies today ranges widely ($10^{-4}-1\\Zsun$), which reflects the large scatter of metallicity of stars formed at high-redshift.   The G-dwarf problem in the solar neighborhood offers an interesting insight for formation of galactic structure. If we apply the prediction of the metallicity distribution to solar neighborhood, we encounter the G-dwarf problem: the prediction of the metallicity distribution does not differ much from that of the closed-box model. In our case, however, the G-dwarf problem would be solved if bulges are early entities and disks are later additions.   We have briefly discussed observational tests for each item of the predictions. In many aspects the predictions are supported by observational evidence, but for some aspects the predictions do not seem to agree with the current interpretation of the observations. The currently available observations, however, are sometimes controversial among authors or do not offer unique interpretations free from various working assumptions. We consider that decisive tests will be a task for the future."
    },
    "0011/astro-ph0011191_arXiv.txt": {
        "abstract": "We have analyzed the observed image separations of the gravitationally lensed images of QSOs for a possible correlation with the source redshift. Contrary to the previously noted anti-correlation based on a smaller data set, no correlation is found for the currently available data. We have calculated the average image separations of the lensed QSOs as a function of source redshifts, for isothermal spheres with cores in a flat universe, taking into account the amplification bias caused by lensing. The shape of the distribution of average image separation as a function of redshift is very robust and is insensitive to most model parameters. Observations are found to be roughly consistent with the theoretical results for models which assume the lens distribution to be (i) Schechter luminosity function which, however, can not produce images with large separation and (ii) the mass condensations in a cold dark matter universe, as given by the Press-Schechter theory if an upper limit of 1-7$\\times 10^{13}$ M$\\odot$ is assumed on the mass of the condensations. ",
        "introduction": "The phenomenon of gravitational lensing is extremely useful for understanding the large scale structure of the universe.  Studying the observed properties of the images in individual lenses can provide knowledge of the mass distribution and mass to light ratio in these lenses (Grogin and Narayan, 1996; Tyson, Kochaski and Dell'Antonio, 1996).  A statistical analysis of the observed lens systems on the other hand can help us to restrict the values of the cosmological parameters (Falco, Kochanek and Munoz, 1998; Link and Pierce, 1998). Based on the frequency of gravitational lensing, upper limits have been placed on the value of the cosmological constant (Fukugita et al 1992, Kochanek, 1996).  It has recently been pointed out by Park and Gott (1997, hereafter PG) that the image separations of the observed lenses are strongly negatively correlated with source redshifts, z$_{\\rm s}$, the anti-correlation being much stronger than that predicted by standard cosmologies.  PG considered several possible ways to strengthen the anti-correlation.  A steeper mass profile, merger of galaxies and increase in their mass by in-fall strengthen the anti-correlation, but not sufficiently to explain the observations.  PG considered point masses and singular isothermal spheres (SISs) for the lenses and assumed all the lenses to be identical. They also, did not consider the effect of amplification bias in their calculation of average image separation.  Williams (1997, hereafter LLRW) showed that the theoretical upper limit to the image separations shows a strong anti-correlation with the source redshift, consistent with the observations, provided (i) the lensing galaxies have logarithmic surface mass densities that gradually change with radius, (ii) there is a dispersion in the lensing properties of galaxies like the central surface mass density or velocity dispersion and (iii) the characteristic length scale of dark matter halos of galaxies scales as L$^a$, with a$\\simeq$ 0.4.  More data on lensed QSOs have become available since 1997 and the number of lensed QSOs is now roughly twice the number in the PG sample.  The PG sample also contained some doubtful lenses. The new data set has to be examined afresh for the presence or absence of the correlation between image separation and source redshift. The main aim of the present paper is to analyze the currently available data to look for a possible correlation. We then want to compare the observed distribution of average image separation as a function of source redshift with the results of a detailed calculation of statistical lensing of galaxies, with a view of obtaining constraints on the values of various parameters entering the calculations. The plan of the paper is as follows. In section 2 we present the analysis of the data using various statistical tests to quantify the presence of any correlation. Calculations for lensing statistics in particular of the average image separation as a function of source redshift, for flat universe, for realistic lens distributions, taking into account the amplification bias, for different cosmological models are presented in section 3. In section 4 and 5 we present the results and conclusions respectively. ",
        "conclusions": "We have analyzed the observed image separations of lensed QSOs for a possible correlation with the source redshift. A correlation was earlier noted for the then available data by PG. The present data of 39 confirmed or likely lensed QSOs, even when combined with 5 wide separation QSO pairs which are doubtful lenses, do not show any statistically significant correlation. LLRW had shown that if the core radii scale with luminosity as L$^{0.4}$ then an anti-correlation is expected theoretically because of the presence of dispersion in lensing properties of galaxies . We have calculated the average image separations of the lensed QSOs as a function of source redshifts, for isothermal spheres with cores in a flat universe, taking into account the amplification bias caused by lensing. We do not find the strong anti-correlation stipulated by LLRW. In fact the shape of the distribution of average image separation as a function of source redshift is very robust and is insensitive to the change in parameters. As a result we are unable to obtain meaningful constraints on any of the parameters. The assumption of the Schechter luminosity function is unable to produce separations larger than 6 $\\arcsec$. The use of PS distribution on the other hand yields large number of wide separation lenses even for non-zero core radius and nessecitates the assumption of a cutoff on the mass of the condensations or large values of core radii for condensations with large masses. It may be noted that the PS distribution has also been found to be remarkably successful in explaining the observed distributions of the QSO absorption lines (Das \\& Khare, 1999)."
    },
    "0011/astro-ph0011158_arXiv.txt": {
        "abstract": "We model the intermediate time evolution of a ``jetted'' gamma-ray burst by two blobs of matter colliding with the interstellar medium. We follow the hydrodynamical evolution of this system numerically and calculate the bremsstrahlung and synchrotron images of the remnant. We find that for a burst energy of $10^{51}\\unts{erg}$ the remnant becomes spherical after $\\sim 5000$ years when it collects $\\sim 50M_\\odot$ of interstellar mass. This result is independent of the exact details of the GRB, such as the opening angle. After this time a gamma-ray burst remnant has an expanding sphere morphology. The similarity to a supernova remnant makes it difficult distinguish between the two at this stage.  The expected number of non-spherical gamma-ray burst remnants is $\\sim0.05$ per galaxy for a beaming factor of $0.01$ and a burst energy of $10^{51}\\unts{erg}$. Our results suggest that that the double-shell object \\dem~is not a GRB remnant. ",
        "introduction": "\\label{sec:Intro}  If a $\\gamma$-ray burst (GRB) originate within a galactic disk then the large deposition of energy will result in a blast wave whose initial phase produces the afterglow. The late phase of the blast wave evolution would result, as noted by \\cite{chevalier74} in the context of supernova remnants (SNRs), in a cool expanding \\HI~shell. The shell will remain distinct from its surrounding until it has slowed down to a velocity of $\\approx 10\\unts{km\\,s^{-1}}$. This would happen within $2.3E_{51}^{0.32}10^6\\unts{yr}$ where $E_{51}$ is the initial energy in units of $10^{51}\\unts{erg}$ \\citep{loeb98}.  The current rate of GRBs is one per $\\sim10^7\\unts{yr}$ per galaxy \\citep{schmidt99}. This leads us to expect a few remnants per galaxy at any given time. However SN are $10^5$ times more frequent and it would be difficult to distinguish between a SNR and a GRB remnant.  Today there is mounting evidence that some GRBs are beamed \\citep{halpern99,harrison99,kulkarni99,sari99a,kuulkers00,salmonson00}. Beamed GRBs will illuminate only a fraction $f_b$ of the sky, and their rate should be higher by a factor of $f_b^{-1}$. With $f_b\\sim 0.01$ we would expect a hundred GRB remnants per galaxy. At an early stage the morphology of a ``jetted'' GRB remnant would be very different from a spherical explosion. This can be used to identify those remnants. We study this phase here.  Both beamed GRBs and SN deposit a comparable ($\\sim 10^{51}\\unts{erg}$) energy into the ISM. In both cases the evolution is expected to be similar since both are in the Sedov \\citep{Sedov59} regime where all the energy is the initial explosion energy and all the mass is in the surrounding ISM. A key distinguishing feature unique to GRB remnants could be their beamed nature which we expect would lead to a distinct double shell morphology at intermediate times (the late time behavior of the GRB remnant is expected to be spherical in any case). In order to establish how many \\HI~shells are GRB remnants we must find out the expected morphology of GRB remnants and how long they remain non-spherical, and distinguishable from SNRs.  Expanding \\HI~shells have been found in many spiral galaxies \\citep{tenorio-tagle88}. The interior of these shells is relatively empty and their current expansion velocity is in the order of tens of Km s$^{-1}$. Models for the origin of these \\HI~supershels involve a large number of spatially correlated supernova events \\citep{heiles79,heiles84}, infall of massive gas clouds on to the galactic plane \\citep{tenorio-tagle88} and flaring of radio lobes formed by jets ejected from the galactic nucleus during an active phase \\citep{gopal-krishna00}. It has been previously noted \\citep{efremov98,loeb98} that a subset of these \\HI~supershells may be the late signature left by GRBs on the interstellar medium (ISM). Establishing how many of the \\HI~shells are GRB remnants would make it possible to directly estimate the local rate of GRBs, determine $\\epsilon$, the efficiency of converting the explosion energy into $\\gamma$-rays, and the beaming factor $f_b$ \\citep{loeb98}.  We model the intermediate evolution of a beamed GRB by two blobs of dense material moving into the ISM in opposite directions and we follow numerically their hydrodynamical evolution. Our results can be rescaled to fit a variety of initial energies. We find that at a time of $\\sim 5000\\unts{yr}$, when the ratio $Mc^2/E_0$ between the accumulated mass, $M$, and the initial GRB energy,$E_0$, reaches $\\sim 9.6\\times10^4$ the remnant becomes spherical, similar in shape to a SNR. This value is independent of the exact initial details of our model such as the opening angle, the velocity and the morphology. We compare our results to the \\HI~shells \\dem~\\citep{Williamsetal97} previously classified as two colliding SNRs. As the accumulated mass there is much larger it is most likely not a GRB remnant.  We describe our model and the numerical methods in section 2 . In section 3 we describe our results. Discussion and summary including a comparison with \\dem~are given in section 4. ",
        "conclusions": "We follow the hydrodynamic evolution of two blobs colliding with the ISM. This scenario is a model for the intermediate time behavior of matter ejected from a central engine producing a GRB. Our results can be rescaled to fit a variety of initial energies and ISM densities. The two shells touch and a shock forms along the equatorial plane between them at $t\\sim 50-160(E_{51}/\\ri)^{-1/3}\\unts{yr}$ when $\\mu = 1-5\\times 10^3$. We show that the late time remnant is insensitive to the exact initial morphology, angular width and density of the ejecta. Although initially the remnant may be highly non-spherical, the ratio between its height and radius will approach unity and it will eventually become spherical in shape after a time of $\\sim 5\\times 10^3E_{51}^{1/3}\\ri^{-1/3} \\unts{yr}$ when $\\mu\\gtrsim 9.6\\times 10^4$. After this time it will be difficult to distinguish a GRB remnant from a SNR using the morphology.  The expected number of non-spherical GRB remnants is, therefore, $5\\times10^{-4}f_b^{-1}E_{51}^{1/3}\\ri^{-1/3}$ per galaxy (using the observed present GRB rate \\citep{schmidt99} of $10^{-7}\\unts{yr^{-1}\\,gal^{-1}}$).  The results of \\cite{gaensler98} show a tendency for the bilateral axis of the non-spherical SNRs to be aligned with the galactic plane. This presents strong evidence in favor of an extrinsic model for the origin of non-spherical SNRs. Using our results, we can propose an alternative explanation to the origin of some of the highly non-spherical SNRs and \\HI~ shells. Instead of assuming an {\\em extrinsic} model, namely spherical energy deposition into a non isotropic medium we propose an {\\em intrinsic} model: non spherical energy deposition into an isotropic medium. Our GRB remnant model can explain non-spherical SNRs with a two shell morphology provided that $\\mu\\lesssim 9.6\\times 10^4$. One extreme and well studied case, the non-spherical SNR \\dem~\\citep{Williamsetal97}, has a distinct double shell morphology. The external model clearly fails here. However the ratio $\\mu$ measured for \\dem~is $\\gtrsim 7.1\\times10^5$ and a a GRB remnant would already be spherical at this stage, suggesting that \\dem~ is not a GRB remnant.  This conclusion could be revised if for some reason the initial blobs stay confined for a much longer period (so that eq. 1 is not satisfied) or if the solution becomes radiative before becoming spherical and the radiative cooling slow down the evolution."
    },
    "0011/astro-ph0011534_arXiv.txt": {
        "abstract": "I study the formation of Comptonization spectra in spherically symmetric, fast moving media in a flat spacetime. I analyze the mathematical character of the moments of the transfer equation in the system-frame and describe a numerical method that provides fast solutions of the time-independent radiative transfer problem that are accurate in both the diffusion and free-streaming regimes. I show that even if the flows are mildly relativistic ($V\\sim 0.1$, where $V$ is the electron bulk velocity in units of the speed of light), terms that are second-order in $V$ alter the emerging spectrum both quantitatively and qualitatively. In particular, terms that are second-order in $V$ produce power-law spectral tails, which are the dominant feature at high energies, and therefore cannot be neglected. I further show that photons from a static source are upscattered by the bulk motion of the medium even if the velocity field does not converge. Finally, I discuss these results in the context of radial accretion onto and outflows from compact objects. ",
        "introduction": "Compton scattering in quasi-spherical accretion flows is thought to be responsible for the X-ray spectra of many accreting compact objects. For example, the hard X-ray spectra of weakly-magnetic accreting neutron stars can be accounted for if a geometrically thick, hot medium surrounds the neutron star and inner accretion disk (see, e.g., Psaltis, Lamb, \\& Miller 1995). The spectral signature of isolated neutron stars accreting from the interstellar medium, which is required for their positive identification, is determined mainly by the efficiency of Compton upscattering of thermal photons emitted from the stellar surfaces (see, e.g., Zampieri et al.\\ 1995). Moreover, Compton scattering plays a major role in the formation of the high-energy spectra of advection-dominated accretion flows, which may be the relevant mode of accretion onto galactic and supermassive black holes (see, e.g., Esin et al.\\ 1997).  Compton scattering in static and moving media, even in plane-parallel or spherical symmetry, couples strongly all three independent phase-space coordinates of the photons, namely their energy, direction of propagation, and spatial coordinate. For this reason, Monte Carlo methods have been used extensively in calculating Comptonization spectra, because of their easy implementation and natural treatment of the diverse length scales in an accretion flow (see, e.g., Laurent \\& Titarchuk 1999 for a recent study; methods and results based on Monte Carlo treatments have been reviewed by Pozdnyakov, Sobol, \\& Sunyaev 1983 and will not be discussed here).  Alternatively, Compton scattering can be studied by solving the energy-dependent radiative transfer equation or its moments. Implementing such a method is more difficult and often relies on a number of approximations for describing the interaction of radiation with matter (see, however, Poutanen \\& Svensson 1996 and Hsu \\& Blaes 1998 for complete solutions in static media). However, at the same time it is substantially faster and offers significant advantages when the radiative transfer problem is coupled to a solution of the fluid dynamics, as well as in treating steep spectra and processes that do not conserve photons.  Over the last twenty five years several authors have derived the photon kinetic or radiative transfer equation and their moments for Compton scattering in static and moving media (see, e.g., Kompaneets 1957; Babuel-Peyrissac \\& Rouvillois 1969; Pomraning 1973; Chan \\& Jones 1975; Payne 1980; Blandford \\& Payne 1981a; Thorne 1981; Fukue, Kato, \\& Matsumoto 1985; Madej 1989, 1991; Titarchuk 1994). The moment equations derived by these various sets of authors have been widely used in studies of Comptonization by static media (see, e.g., Katz 1976; Shapiro, Lightman, \\& Eardly 1976; Sunyaev \\& Titarchuk 1980) or by strong shocks and accretion flows onto compact objects (see, e.g., Blandford \\& Payne 1981b; Payne \\& Blandford 1981; Lyubarskij \\& Sunyaev 1982; Colpi 1988; Riffert 1988; Mastichiadis \\& Kylafis 1992; Titarchuk \\& Lyubarskij 1995; Turolla et al. 1996; Titarchuk, Mastichiadis, \\& Kylafis 1997), either under the diffusion approximation or with closure relations that were specified a priori. Most of the above analyses of Compton scattering in moving media were performed in the system frame and to first order in the electron bulk velocity. On the other hand, Schmid-Burgk (1978), Zane et al.\\ (1996), and Titarchuk \\& Zannias (1998) have solved the general-relativistic radiative transfer equation for Compton scattering without the need for specifying a priori closure relations.  In Paper I (Psaltis \\& Lamb 1997) of this series, the time-dependent photon kinetic and radiative transfer equations were derived, as well as their zeroth and first moments that describe absorption, emission, and spontaneous and induced Compton scattering in static and moving media, correcting various errors in the literature. The system-frame equations that were derived are valid to first order in $\\epsilon/m_e$ and $T_e/m_e$, and to {\\it second} order in $V$, where $m_e$ and $T_e$ are the electron mass and temperature, $\\epsilon$ is the photon energy, and $V$ is the bulk velocity of the electrons in units of the speed of light; the fluid-frame equations that were derived are valid for arbitrary values of the bulk velocity $V$. Using these equations it was argued in Paper I that the effects of Comptonization by the bulk electron velocity that are described by the terms that are second-order in $V$ are qualitatively different than and can become at least as important as the effects described by the terms that are first-order in $V$, even when $V$ is small.  As a result, these second-order terms should generally be retained (see also Yin \\& Miller 1995).  In this paper, I describe a numerical algorithm for the solution of time-independent radiative transfer problems in systems with spherical symmetry.  Even though the natural reference frame for solving transport problems is the one comoving with the flow, I work in the system frame, which can be any inertial frame in which the central object is at rest. This is important for the fast algorithm described below, because the moment equations in the system frame are always parabolic (as shown in \\S3) in contrast to the fluid-frame equations (Turolla, Zampieri, \\& Nobili 1995; Smit, Cernohorsky, \\& Dullemond 1997; see also K\\\"orner \\& Janka 1992). Moreover, in problems of mass accretion onto a compact star, the boundary conditions are often more easily implemented in the system frame.  I also neglect any general relativistic effects, which can be shown to introduce only small quantitative corrections to the emerging spectra (Papathanassiou \\& Psaltis 2000), and truncate the transfer equation keeping only terms up to second order in $V$.  The algorithm used is a generalization of the method of variable Eddington factors (Mihalas 1980; see also Mihalas 1978, p.~157--158). It is therefore accurate for systems of arbitrary optical depth and does not require a priori specification of closure relations. It is based on the iterative solution of both the transfer equation and the systems of its first two moments and hence the validity and accuracy of the solution can be verified explicitly at the end of each run. The method also requires only a small number of iterations in order to converge to the solution and is therefore ideal for future extensions to time-dependent and multi-dimensional transport problems.  In \\S2, I summarize the results from Paper I and introduce the notation. In \\S3, I discuss the mathematical character of the system of equations and in \\S4, I present the numerical algorithm. In \\S5, I describe the results of numerical solutions of the transfer equations for a variety of situation. Finally, in \\S6, I briefly summarize the implications of the above for Comptonization spectra in static and moving media. ",
        "conclusions": "In the previous sections I described a method for solving numerically the radiative transfer equation that describes Compton scattering in spherically symmetric, static and moving media, and applied it to a variety of cases related to flows around compact objects. I showed that the first Eddington factor, which describes the degree of isotropy of the radiation field at a given point in space, may depend on photon energy, even though the photon mean-free-path is energy independent, if the photons are preferably upscattered or downscattered in energy at each scattering. I also demonstrated that the terms that are first-order in the electron bulk velocity alter the input spectrum in a way that is qualitatively and quantitatively different than the effect of the terms that are of second-order.  The systematic upscattering of photons caused by bulk Comptonization is described by the terms that are second-order in $V$. This becomes apparent by the fact that, at each scattering with a moving electron, the average fractional energy change of a photon in the system frame is always positive, independent of the photon energy, and equal to $4V^2/3$ (see, e.g., Rybicki \\& Lightman 1979, ch.~7; see also the discussion in Psaltis \\& Lamb 1997). If the distribution of photon residence times in the Comptonizing medium drops exponentially to zero at late times, then a power-law tail is produced at high photon energies (for total scattering optical depths $\\gtrsim 1$), in analogy to the generation of power-law tails by thermal Comptonization.  The velocity at which the terms that are second-order in $V$ become important depends on the ratio of the photon mean-free-path $\\lambda_{\\rm mfp}$ to the shortest characteristic length scale $L$ in the problem. In general, when $V(L/\\lambda_{\\rm mfp}) \\sim 1$, i.e., in regions close to and within the photon-trapping radius, the terms that are first- and second-order in $V$ become comparable in size (Yin \\& Miller 1995; Psaltis \\& Lamb 1997a) and hence the latter should not be neglected.  The generalization of the method to time-independent multi-dimensional configurations is conceptually easy. However, a problem with axial symmetry requires the solution of the radiative transfer equation in five dimensions in phase space and the introduction of a large number of independent variable Eddington factors. Moreover, the presence of high flow velocities in multi-dimensions introduces sharp gradients in the variable Eddington factors, which cannot be guessed a priori, as required for the first iteration (Dykema, Klein, \\& Castor 1996). Even with these problems, however, the method of variable Eddington factors will still offer a significant advantage, when compared to any other iterative method, which arises from the fact that the Eddington factors are bound to lie over a very narrow range of values.  The fast convergence of the numerical algorithm presented here makes it also ideal for future extensions to transport problems coupled to time-dependent hydrodynamics. For the case of accretion onto and outflows from compact objects, the characteristic photon diffusion timescale from a flow of dimension $R$ is of order \\begin{equation} t_{\\rm dif}\\simeq\\frac{R\\tau}{c}=30\\left(\\frac{R}{10^6~\\mbox{cm}}\\right) \\left(\\frac{\\tau}{1}\\right)~\\mu\\mbox{s}\\;, \\end{equation} which is significantly smaller than the corresponding $t_{\\rm dyn}\\sim 1$~ms dynamical timescale. As a result, for most cases the structure of the radiation field can be obtained by solving the time-independent radiative transfer problem for each snapshot of the fluid properties. Given that the variable Eddington factors will be only marginally different between successive timesteps, the convergence of the algorithm will be even faster than what is shown in Figure~2."
    },
    "0011/astro-ph0011252_arXiv.txt": {
        "abstract": "I review the reason for considering the prime purpose of the program of measurements of the fundamental parameters of cosmology to be the tests of cosmological models. I comment on the philosophy by which we are approaching this goal, offer an assessment of where we stand, and present some thoughts on where the tests may be headed. ",
        "introduction": "These Proceedings document impressive progress toward a satisfactory completion of the great program of measurements of the parameters of cosmology that commenced in the 1930s. It has taken a long time, and has brought into play many phenomena and measurements that could not have been anticipated in the 1930s. We may at last be approaching closure of this program, and it is appropriate to reflect on why we are so interested in these measurements and how it informs our interpretation of the results. ",
        "conclusions": "The evidence assembled in these Proceedings favors the existence of several kinds of matter: one that acts like Einstein's cosmological constant, with density parameter $\\Omega _\\Lambda\\sim 0.75$, nonbaryonic low pressure matter with density parameter $\\Omega _{\\rm DM}\\sim 0.2$, baryons with density parameter $\\Omega _{\\rm baryons}\\sim 0.05$, and neutrinos with $\\Omega _\\nu\\sim 0.001$. I believe it is too soon to add the first number to the list of firmly established elements of the science of cosmology, because it depends on the model for structure formation, and some of us see apparent problems with the model. People have been discussing the cosmological parameters for seven decades; we can wait a few more years to determine  whether we have got the science right.  If advances in the applications of the cosmological tests firmly established the values of the fundamental parameters of the Friedmann-Lema\\^\\i tre model it would mean general relativity theory has satisfied demanding tests on the scales of cosmology, and that we have a well-tested history of structure formation. But I would not be surprised to find that this advance leaves us with the challenge of establishing the physics of the very early universe."
    },
    "0011/astro-ph0011064_arXiv.txt": {
        "abstract": "We have constructed models for the sun at three stages of its evolution: a zero-age main sequence model, the present sun, and a subgiant model.  For each model, the turbulent pressure and turbulent kinetic energy were calculated from 3-d radiative hydrodynamical simulations (described in the poster by Robinson et al.), and inserted into the 1-d stellar models.  We note that in these simulations, the turbulent pressure is not a free parameter, but can be computed from the resulting velocity field.  We show the calculated p-mode frequencies for the model of the present sun, with and without turbulent pressure, and compare them to the observed solar frequencies. When the turbulent pressure is included in the models, the calculated frequencies are brought closer to the observed frequencies in the sun by up to two $\\mu Hz$, strictly from structural effects.  The effect of including turbulent pressure on p-mode frequencies is also shown for the zero-age main sequence model.  Our models also suggest that the importance of turbulent pressure increases as the star evolves into the subgiant region.  We discuss the importance of also including realistic turbulence as well as radiation in the non-adiabatic calculation of oscillation frequencies. ",
        "introduction": "The purpose of this paper is to calculate the effect of including the turbulent pressure $P_{turb}$, derived from the 3-d radiative hydrodynamic simulations of Robinson et al. (described in a companion poster paper), into the hydrostatic structure of the outer layers of the models, where $P_{turb}$ is known to be significant; and to determine the frequency shift caused by the inclusion of $P_{turb}$ on the p-mode oscillation frequencies.  The turbulent pressure was calculated for three models of the evolving sun, corresponding to the sun on the zero-age main sequence (ZAMS), the present sun, and the future sun as a subgiant. ",
        "conclusions": ""
    },
    "0011/astro-ph0011408_arXiv.txt": {
        "abstract": "We study the accretion efficiency of neutron stars in long--period binaries ($P \\ga 20^{\\rm  d}$) which accrete from a giant companion. Using $\\alpha$--disc models and taking into account the effect of irradiation of the accretion disc by the central accretion light source we derive explicit expressions for the duty cycle $d$ and the accretion efficiency $\\eta$ in terms of the parameters of the binary system and the disc instability limit cycle. We show that the absence of millisecond pulsars in wide binaries with circular orbits and periods $P \\ga 200^{\\rm  d}$ can be understood as a consequence of the disc instability if the duration of the quiescent phase between two subsequent outbursts is at least a few decades. ",
        "introduction": "Ever since the detection of the first millisecond pulsar (Backer et al. 1982) it has been clear that a ms-pulsar is a neutron star (NS) which has been spun up by accretion in a close binary system (for a review see Bhattacharya \\& van den Heuvel 1991). Because the amount of mass needed to spin up a NS to ms periods is rather small, i.e. $0.05 \\msun \\la \\Delta M_{\\rm  NS} \\la 0.1 \\msun$ (Burderi et al. 1999) and the mass available from the donor is typically many times larger, i.e. typically of order $1 \\msun$, it appeared as if there was no problem spinning up a NS to such short periods. Yet in recent years, evidence has accumulated which shows that NS are inefficient accretors. Thorsett \\& Chakrabarty (1999) noted that in all cases where the mass of a NS could be measured {\\it with precision} it turned out to be in the very narrow interval $1.3 \\msun \\la M_{\\rm NS} \\la 1.4 \\msun$. Even if one makes the extreme assumption that all NSs are born with exactly the same mass, this would imply that none of these NSs could have accreted more than $\\sim 0.1 \\msun$, despite the fact that some of these NSs have been spun up by accretion and the mass available from the donor star was probably much larger than $\\sim 0.1 \\msun$. In another context, King \\& Ritter (1999) have presented further evidence that NSs can be very inefficient accretors. The main reason for this low accretion efficiency is the frequent disparity between the mass transfer rate provided by the donor star and the maximum rate at which the NS can accrete. The latter is approximately the Eddington rate $\\dot M_{\\rm E} \\approx 1.5 \\times 10^{-5} \\msun {\\rm yr}^{-1}$, whereas, depending on circumstances, the former can be larger by many orders of magnitude.  One of the obvious places in which NSs can be spun up to ms--pulsars are wide binaries in which a NS accretes from a giant. Tauris \\& Savonije (1999) have recently devoted a detailed study to this formation channel. One of their predictions is that, in principle, ms--pulsars should form in binaries with final orbital periods of up to $P \\simeq 1000^{\\rm d}$, even when taking into account that the mass transfer rate in long--period systems ($P \\ga 20^{\\rm d}$) exceeds the Eddington accretion rate. In addition, their calculations predict a correlation between the final orbital period and the mass of the ms--pulsar. Unfortunately, this cannot be tested at present. But their first prediction can. And here observations show, contrary to their prediction, that there are no ms--pulsars in wide binaries with a circular orbit and a period $P \\ga 200^{\\rm d}$ (see e.g. table 1 in Taam, King \\& Ritter 2000). In this context, Li \\& Wang (1998) have already noted that the accretion efficiency of the NS in such a binary is far lower than has been assumed by Tauris \\& Savonije (1999) if one takes into account the fact that in all these wide binaries the accretion disc around the NS is thermally unstable and accretion is transient (King et al. 1997). This is because a) during an outburst the mass flow rate through the disc is much larger than the mass transfer rate from the donor, which, in turn, is already super--Eddington if $P \\ga 20^{\\rm d}$, and b) because a NS cannot accrete in excess of the Eddington rate. In fact, Li \\& Wang (1998) found that if the duty cycle of the disc instability limit cycle is $d \\la 10^{-2}$ then in systems with a final orbital period $P_{\\rm f} \\ga 100^{\\rm d}$ the total mass accreted by the NS is $\\Delta M_{\\rm  NS} \\la$ few $10^{-2}\\msun$, too little to spin up a NS to ms periods.  In this paper we elaborate on the idea of Li \\& Wang (1998) and derive explicit expressions for the duty cycle and the accretion efficiency of the NS in terms of the parameters of the binary system and the outburst cycle. Our calculations show that if the duration of the quiescent phase between two subsequent outbursts exceeds a few decades, the accretion efficiency becomes so small that in binaries with a final orbital period $P_{\\rm f} \\ga 200^{\\rm d}$, the NS cannot accrete enough mass to become a ms pulsar. ",
        "conclusions": "Li \\& Wang (1998) have shown that the accretion efficiency of neutron stars which accrete from a giant is very small and, in fact, becomes too small for spinning up the neutron star to millisecond spin periods in systems with a final orbital period $P_{\\rm f} \\ga 100^{\\rm d}$ if the duty cycle $d \\la 10^{-2}$. Here we have worked out explicit expressions for $d$ and the accretion efficiency $\\eta$ in terms of the  binary parameters and the quiescent time $t_{\\rm q}$ of the outburst cycle in the framework of an irradiated $\\alpha$--disc model. Our calculations show that both, $d$ and $\\eta$ are small if $P > P_{\\rm crit}$ and the mass flow rate at the onset of an outburst is highly super--Eddington. Whether or not $d$ and $\\eta$ are really small enough to prevent neutron stars in systems with $P_{\\rm f} \\ga 200^{\\rm d}$ being spun up to ms--periods depends crucially on the duration of quiescence $t_{\\rm q}$. If $t_{\\rm q} \\ga$ many decades, as some observations of BHSXTs suggest, then both  $\\eta$ and $d$ are very small indeed. Unfortunately, determining $t_{\\rm q}$ in the framework of our simple model is not possible. For this fully time--dependent calculations of the disc instability limit cycle would be needed.  However there is another way that we can get a lower limit on $d$. We know that pulsars in wide circular binaries must have descended from SXTs. Yet we do not know of a single transient with a period longer than 11.8$^{\\rm d}$ (GRO J1744--28), even though X--ray satellites would probably have detected an outburst from such a system anywhere in the Galaxy within the last $\\sim 30$~yr. This must mean that outbursts are rather infrequent, i.e. $t_{\\rm q}$ is long. We can make this more precise by defining the following quantities:  Let $n$ be the frequency of X--ray outbursts from all systems in the Galaxy, $N_{\\rm PSR,obs}$ the currently known number of the pulsar binaries in question, $N_{\\rm PSR}$ the current number of the pulsar binaries in question in the Galaxy, $\\tau_{\\rm prog}$ the lifetime of a progenitor system in the phase of nuclear time scale mass transfer, and $\\tau_{\\rm PSR}$ the lifetime of the pulsar in the pulsar binaries in question. Here $\\tau_{\\rm prog}$ and $\\tau_{\\rm PSR}$ are suitable averages over the relevant population.  Now, the inventory of the binary pulsars in question is incomplete because of flux limitation, dispersion and beaming, whereas an X--ray outburst in one of the progenitor systems would probably be seen wherever in the Galaxy the outburst occurs. We correct for this incompletness by introducing a `filling factor' $f < 1$. In a stationary situation we then must have  \\begin{equation} t_{\\rm q} = {N_{\\rm PSR} \\over n}\\; {\\tau_{\\rm prog}\\over\\tau_{\\rm PSR}} = {N_{\\rm PSR,obs}\\over n\\;f}\\; {\\tau_{\\rm prog}\\over \\tau_{\\rm PSR}}. \\end{equation} $\\tau_{\\rm PSR}$ must be of the order of the spin-down time scale $P/2\\dot P \\approx 10^8 ~{\\rm yr}$, $\\tau_{\\rm prog}$, on the other hand, is of the order of the nuclear time scale $t_{\\infty}$ defined in Ritter (1999, Eq. 14) which, for the systems in question, i.e. those which end their evolution with a long orbital period $P_{\\rm orb} \\ga 200^{\\rm d}$, is of the order $10^8~{\\rm yr}$. With $N_{\\rm PSR,obs}=2$ (i.e. PSR B0820+02 and PSR J1803-2712), $f \\approx 0.2$ and from the fact that we have not seen a single outburst in the past 30~{\\rm yr} this yields the lower limit $t_{\\rm q} > (60~{\\rm yr})/f \\ga 300~{\\rm yr}$. We would also like to emphasize that we have adopted a very conservative value for $f$. The real value is likely to be smaller still and thus the estimate for $t_{\\rm q}$ even larger because already the canonical beaming factor for pulsars is of order $0.2$.  Thus it is very likely that the lack of ms--pulsars in binaries with white dwarf companions and $P_{\\rm orb} \\ga 200^{\\rm d}$ is entirely a consequence of dwarf nova--like disc instabilities in long--period binaries.  Finally, we note that spinning up a NS in a long--period binary faces yet another obstacle: during the long quiescent phases the NS acts as a propeller (Illarionov \\& Sunyaev 1975) and thereby is spun {\\em down}, making the spin-up to ms spin periods even more difficult."
    },
    "0011/astro-ph0011122_arXiv.txt": {
        "abstract": "The fragmentation of molecular clouds (MC) into proto--stellar cores is a central aspect of the process of star formation. Because of the turbulent nature of super--sonic motions in MCs, it has been suggested that dense structures such as filaments and clumps are formed by shocks in a turbulent flow. In this work we present strong evidence in favor of the turbulent origin of the fragmentation of MCs.  The most generic result of turbulent fragmentation is that dense post shock gas traces a gas component with a smaller velocity dispersion than lower density gas, since shocks correspond to regions of converging flows, where the kinetic energy of the turbulent motion is dissipated.  Using synthetic maps of spectra of molecular transitions, computed from the results of numerical simulations of super--sonic turbulence, we show that the dependence of velocity dispersion on gas density generates an observable relation between the rms velocity centroid and the integrated intensity (column density), $\\sigma(V_0)-I$, which is indeed found in the observational data. The comparison between the theoretical model (maps of synthetic $^{13}$CO spectra), with $^{13}$CO maps from the Perseus, Rosette and Taurus MC complexes, shows excellent agreement in the $\\sigma(V_0)-I$ relation.  The $\\sigma(V_0)-I$ relation of different observational maps with the same total rms velocity are remarkably similar, which is a strong indication of their origin from a very general property of the fluid equations, such as the turbulent fragmentation process. ",
        "introduction": "The importance of turbulence in the process of star formation was recognized long ago (Von Weizs{\\\"a}cker 1951), and was discussed in a seminal paper \\nocite{von51} by Larson (1981). Several successive works have tried to use the observational data to relate different properties of MCs with the physics of laboratory and numerical turbulence, such as power spectra of kinetic energy \\footnote{See Leung, Kutner \\& Mead 1982; Myers 1983; Quiroga 1983;  Sanders, Scoville \\& Solomon 1985;   Goldsmith \\& Arquilla 1985; Dame et al.\\ 1986; Falgarone \\& P\\'{e}rault 1987.} \\nocite{Leung+82,Myers83,Quiroga83} \\nocite{Sanders+85} \\nocite{Goldsmith+Arquilla85,Dame+86,Falgarone+Perault87}, probability distributions of velocity and velocity differences \\footnote{ See Scalo 1984; Kleiner \\& Dickman 1985, 1987; Hobson 1992; Miesch \\& Bally 1994; Miesch \\& Scalo 1995; Lis et al. 1996; Miesch, Scalo \\& Bally 1999.}, intermittency (Falgarone \\& Phillips 1990; Falgarone \\& Puget 1995; Falgarone, Pineau Des Forets, \\& Roueff 1995), \\nocite{Falgarone+Phillips90} \\nocite{Falgarone+Puget95} \\nocite{Falgarone+95} and self--similarity \\footnote{See Beech 1987; Bazell \\& D\\'{e}sert 1988, Scalo 1990; Dickman, Horvath \\& Margulis 1990; Falgarone, Phillips \\& Walker 1991; Zimmermann, Stutzki \\& Winnewisser 1992; Henriksen 1991; Hetem \\& Lepine 1993; Vogelaar \\& Wakker 1994; Elmegreen \\& Falgarone 1996.}. \\nocite{Beech87}  \\nocite{Bazell+Desert88}  \\nocite{Scalo90} \\nocite{Dickman+90}  \\nocite{Falgarone+91} \\nocite{Zimmermann+92} \\nocite{Henriksen91} \\nocite{Hetem+Lepine93} \\nocite{Vogelaar+Wakker94} \\nocite{Elmegreen+Falgarone96}  During the last decade, numerical simulations of transonic turbulence (Passot \\& Pouquet 1987; Passot, Pouquet \\& Woodward 1988; L\\'{e}orat, Passot \\& Pouquet 1990; Lee, Lele \\& Moin 1991 \\nocite{Lee+91}; Porter, Pouquet \\& Woodward 1992 \\nocite{Porter+94}; Kimura \\& Tosa 1993; Porter, Woodward \\& Pouquet 1994; Vazquez--Semadeni 1994; Passot, Vazquez--Semadeni \\& Pouquet 1995) and highly super--sonic turbulence (Passot, V\\'{a}zquez--Semadeni, Pouquet 1995; V\\'{a}zquez--Semadeni, Passot \\& Pouquet 1996; Padoan \\& Nordlund 1997, 1999; Stone, Ostriker \\& Gammie 1998; MacLow et al. 1998; Ostriker, Gammie \\& Stone 1999; Padoan, Zweibel \\& Nordlund 2000; Klessen, Heitsch \\& MacLow 2000), on relatively high--resolution numerical grids, have become available, and very detailed comparisons between observational data with turbulence models of MCs have been performed (Padoan, Jones \\& Nordlund 1997; Padoan et al. 1998; Padoan \\& Nordlund 1999; Padoan et al. 1999; Rosolowsky et al. 1999; Padoan, Rosolowsky \\& Goodman 2000).  Recent numerical studies of super--sonic magneto--hydrodynamic (MHD) turbulence have brought new understanding of the physics of turbulence. The most important results are:  \\nocite{MacLow_Puebla98} \\nocite{Stone+98} \\nocite{Kimura+Tosa93} \\nocite{Lee+91} \\nocite{Vazquez-Semadeni+96} \\nocite{Klessen+99} \\nocite{Padoan+97ext} \\nocite{Padoan+98per} \\nocite{Padoan+Nordlund98MHD} \\nocite{Padoan+98cat} \\nocite{Padoan+Nordlund97MHD} \\nocite{Ostriker+99} \\nocite{Passot+95} \\nocite{MacLow+98} \\nocite{MacLow99} \\nocite{Rosolowsky+99} \\nocite{Padoan+2000SCF}   \\begin{itemize} \\item Super--sonic turbulence decays in approximately one dynamical time, independent of the magnetic field strength (Padoan \\& Nordlund 1997, 1999; MacLow et al. 1998; Stone, Ostriker \\& Gammie 1998; MacLow 1999). \\item The probability distribution of gas density in isothermal turbulence is well approximated by a Log--Normal distribution, whose standard deviation is a function of the rms Mach number of the flow (Vazquez--Semadeni 1994; Padoan 1995; Padoan, Jones \\& Nordlund 1997; Scalo et al. 1998; Passot \\& Vazquez--Semadeni 1998; Nordlund \\& Padoan 1999; Ostriker, Gammie \\& Stone 1999). \\item Super--sonic isothermal turbulence generates a complex system of shocks which fragment the gas very efficiently into high density sheets, filaments, and cores (this is the general result of any numerical simulation of super--sonic turbulence). \\item Super--Alfv\\'{e}nic turbulence provides a good description of the dynamics of MCs, and an explanation for the origin of dense cores with magnetic field strength consistent with Zeeman splitting observations. (Padoan \\& Nordlund 1997, 1999).\\footnote{This particular result is supported almost exclusively by our work. A significant fraction of the astrophysical community still favors the traditional idea that a rather strong magnetic field supports MCs against their gravitational collapse. An example of numerical work that favors the traditional idea of magnetic support is Ostriker, Gammie \\& Stone (1999).} \\end{itemize}  \\nocite{Scalo+98} \\nocite{Vazquez-Semadeni94}  We call ``turbulent fragmentation'' the process of generation of high density structures by turbulent shocks. Since random super--sonic motions are ubiquitous in MCs, turbulent fragmentation cannot be avoided: it is a direct consequence of the observational evidence. In some analytical studies it is tacitly assumed that turbulent fragmentation can be neglected if the kinetic energy of random motions is in rough equipartition with the magnetic energy. This assumption is wrong, because motions along magnetic field lines are unavoidable, and so turbulent fragmentation occurs via super--sonic compressions along the magnetic field lines, as discussed by Gammie \\& Ostriker (1996) and Padoan \\& Nordlund (1997, 1999).  \\nocite{Gammie+Ostriker96}  Numerical simulations of turbulence have been used to discuss the turbulent origin of MC structures by Passot \\& Pouquet (1987). Vazquez--Semadeni, Passot \\& Pouquet (1996) and Ballesteros--Paredes, Hartmann \\& Vazquez--Semadeni (1999) have used two dimensional simulations to argue that MCs are formed by turbulence. Padoan \\& Nordlund (1997, 1999) have shown that super--sonic and super--Alfv\\'{e}nic turbulence can explain the origin of magnetized cores in MCs, including the observed field strength--density ($B-n$) relation (Myers \\& Goodman 1988; Fiebig \\& G\\\"{u}sten 1989; Crutcher 1999). The idea that proto--stellar cores and stars are formed in turbulent shocks has been previously discussed by Elmegreen (1993). In that work, an analysis of the gravitational instability of the cores can be found. Here, we present new strong observational evidence in favor of the turbulent shock origin of proto--stellar cores. Such evidence is based on the fact that dense post shock gas traces a gas component with a smaller velocity dispersion than lower density gas, since it maps regions of converging flows, where the kinetic energy of the turbulent motion is dissipated.  \\nocite{Passot+Pouquet87} \\nocite{Ballesteros+99} \\nocite{Myers+Goodman88} \\nocite{Fiebig+Gusten89} \\nocite{Crutcher99} \\nocite{Elmegreen93}  In \\S 2 and 3, numerical simulations and observational data, used in this work, are briefly described. In \\S 4 we compute the rms flow velocity as a function of the gas density, in simulations of super--sonic turbulence, and show that it decreases for increasing values of the gas density. In \\S 5 we show that such general property generates an observable relation between the rms velocity centroid and the integrated intensity (roughly proportional to the surface density), for the \\jco\\ transition, and in \\S 6 the same relation is found in the observational data. Results are discussed in \\S 7, and conclusions are summarized in \\S 8. ",
        "conclusions": "The origin of proto--stellar cores is a fundamental problem in our understanding of the process of star formation. Models that have been proposed to describe proto-stellar cores are based on i) static or quasi--static equilibrium (e.g. Curry \\& Mckee 2000; Jason \\& Pudritz 2000), ii) thermal instability (e.g. Yoshii \\& Sabano 1980; Gilden 1984; Graziani \\& Black 1987), iii) gravitational instability through ambipolar diffusion (e.g. Basu \\& Mouschovias 1994; Nakamura, Hanawa, \\& Nakano 1995; Indebetouw \\& Zweibel 2000; Ciolek \\& Basu 2000), iv) non--linear Alfv\\'{e}n waves (e.g. Carlberg \\& Pudritz 1990; Elmegreen 1990, 1997, 1999), v) clump collisions (e.g. Gilden 1984; Kimura \\& Tosa 1996), vi) super--sonic turbulence (e.g Elmegreen 1993; Klessen, Burkert and Bate 1998; Klessen, Heitsch, \\& Mac Low 2000; Padoan et al. 2000).  Many detailed comparisons between observational data and models, which support the idea of the turbulent origin of the structure and kinematics of molecular clouds, have been presented in our previous papers (Padoan, Jones \\& Nordlund 1997; Padoan et al. 1998; Padoan et al. 1999; Padoan \\& Nordlund 1999; Padoan, Goodman \\& Rosolowsky 2000). Models of numerical turbulence can be compared with the observations by computing i) synthetic stellar extinction measurements, ii) synthetic spectral maps of molecular transitions, iii) synthetic Zeeman splitting measurements, iv) synthetic polarization maps.  Maps of synthetic spectra contain a lot of information about the kinematic, the structure and the thermal properties of molecular clouds, and can be analyzed with different statistical tools. In this work we have presented a new statistical method to analyze spectral--line maps, which is very useful because it probes directly a very general property of super--sonic turbulence, that is the fact that dense gas traces a gas component with a smaller velocity dispersion than lower density gas. This property arises because the gas density is enhanced in regions where the large scale turbulent flow converges (compressions) and the kinetic energy of the turbulent flow is dissipated by shocks. If local compressions are instead due to local instabilities (e.g. gravitational instability, or gravitational instability mediated by ambipolar diffusion), and the large scale motions are only the consequence of local instabilities (as it should be in a self--consistent picture), gas density increases with the flow velocity dispersion (see for example the model by Indebetouw \\& Zweibel 2000), contrary to the observational evidence presented in this work.  \\nocite{Indebetouw+Zweibel2000}  The main conclusion of this work is that every model for the origin of molecular cloud structure and proto--stellar cores should be tested against the $\\sigma(V_0)-I$ relation. Turbulent fragmentation provides a realistic scenario for the origin of proto--stellar cores, which satisfies this new observational constraint.   \\nocite{Passot+88}"
    },
    "0011/astro-ph0011204_arXiv.txt": {
        "abstract": "We analyze the nuclear stellar dynamics of the SB0 galaxy NGC~1023, utilizing observational data both from the Space Telescope Imaging Spectrograph aboard the Hubble Space Telescope and from the ground. The stellar kinematics measured from these long-slit spectra show rapid rotation ($V \\approx 70$ km s$^{-1}$ at a distance of $0\\farcs1 = 4.9$ pc from the nucleus) and increasing velocity dispersion toward the nucleus (where $\\sigma = 295 \\pm 30$ km s$^{-1}$). We model the observed stellar kinematics assuming an axisymmetric mass distribution with both two and three integrals of motion. Both modeling techniques point to the presence of a central dark compact mass (which presumably is a supermassive black hole) with confidence $> 99$\\%. The isotropic two-integral models yield a best-fitting black hole mass of $(6.0 \\pm 1.4) \\times 10^7 \\ M_{\\odot}$ and mass-to-light ratio ($M/L_V$) of $5.38 \\pm 0.08$, and the goodness-of-fit ($\\chi^2$) is insensitive to reasonable values for the galaxy's inclination. The three-integral models, which non-parametrically fit the observed line-of-sight velocity distribution as a function of position in the galaxy, suggest a black hole mass of $(3.9 \\pm 0.4) \\times 10^7 \\ M_{\\odot}$ and $M/L_V$ of $5.56 \\pm 0.02$ (internal errors), and the edge-on models are vastly superior fits over models at other inclinations. The internal dynamics in NGC~1023 as suggested by our best-fit three-integral model shows that the velocity distribution function at the nucleus is tangentially anisotropic, suggesting the presence of a nuclear stellar disk. The nuclear line of sight velocity distribution has enhanced wings at velocities $\\geq 600$ km~s$^{-1}$ from systemic, suggesting that perhaps we have detected a group of stars very close to the central dark mass. ",
        "introduction": "Supermassive black holes (BHs) are widely suspected to be the central engines of quasars and active galactic nuclei. The fact that the comoving number density of quasars peaks at $z \\sim 2 - 3$ (e.g., Schmidt, Schneider, \\& Gunn~1995) and declines by nearly three orders of magnitude by $z = 0$ implies that nearby galaxies should harbor relatively inactive BHs. Searching for these BHs would provide constraints on the evolutionary history of quasars. Also, the most recent census of nearby BHs reveals a remarkable correlation between BH mass and the stellar velocity dispersion measured within the half-light radius of its host galaxy (Gebhardt et al.~2000a; Ferrarese \\& Merritt 2000). This result will have very interesting implications for the formation of galaxy spheroids and the growth of their central BHs (Kormendy 2000).  Dynamical evidence of dark compact objects of mass $\\sim 10^{6.5} \\ - \\ 10^{9.5} \\ M_{\\odot}$ has been accumulating rapidly in the last decade or so (see reviews in Kormendy \\& Richstone~1995; Richstone et al.~1998). Although this evidence in most cases does not strictly exclude alternatives to a BH, it is very difficult to construct plausible alternative models to BHs for three of these BH candidates (the Galaxy, NGC 4258, and M31; see Maoz~1998 and Kormendy \\& Bender~1999) $-$ thus suggesting that many, or perhaps all, of the dark compact objects that have been identified through such dynamical evidence are indeed BHs.  The Hubble Space Telescope (HST) has greatly enhanced this effort to analyze the nuclear dynamics of nearby galaxies. HST's second generation spectrograph, the Space Telescope Imaging Spectrograph (STIS; see Woodgate et al.~1998) can obtain measurements of the nuclear kinematics of either gas or stars with much more efficiency than the first generation HST spectrographs.  Several teams are actively involved in this research. Our group, a subset of the STIS Investigation Definition Team, has selected a galaxy sample including both some of the best black hole candidates and a few galaxies where the ground-based dynamical evidence for a BH is nonexistent or weak. One such galaxy in the latter category is NGC~1023, which is classified as SB(rs)0$-$. Given this morphological type and that $M_B = -19.9$, the luminosity of the bulge component is $M_{B,bulge} \\approx -19.4$ (see Simien \\& deVaucouleurs~1986). We chose NGC 1023 because its intermediate luminosity is consistent with a rotationally supported figure (Faber et al.~1997), removing ambiguity from the interpretation of the nuclear dynamics.  We adopt a distance of 10.2 Mpc (see Faber et al.~1997); thus $1''$ corresponds to a distance at NGC~1023 of 49 pc. The Galactic extinction is $A_B = 0\\fm262$ (Schlegel et al.~1998). ",
        "conclusions": ""
    },
    "0011/gr-qc0011075_arXiv.txt": {
        "abstract": "We show that the non-linear evolution of long-wavelength perturbations  may be important in a wide class of inflationary scenarios. We develop a solution for the evolution of such nonlinear perturbations which is exact to first order in a gradient expansion. As a first application, we demonstrate that in single field models of inflation there can be no parametric amplification of super-Hubble modes during reheating. We consider the implications of the solution for recent discussions of the back-reaction effect of long wavelength perturbations on the background geometry, give a new derivation of the equation of motion of stochastic inflation, and demonstrate that if the (generalized) slow-rolling condition is not satisfied, then inevitably long wavelength vector modes for gravitational fluctuations will be generated.% ",
        "introduction": "The hypothesis that causally unrelated regions of space-time evolve independently is one of the cornerstones in the development of relativity and one of the basic assumptions of all relativistic theories. As an implication of this hypothesis in cosmology we expect that perturbations with long wavelengths\\footnote{Throughout this paper, we use the terms super-Hubble, infrared and long wavelength interchangeably. All refer to modes with wavelengths larger than the Hubble radius.} evolve independently at different points of space, as if they were independent homogeneous patches of the Universe (see e.g. \\cite{Lindebook} and references therein). This expectation has been explored in the linear theory of cosmological perturbations. In this context, it is possible to derive integral constraints (see e.g. \\cite{Traschen}) which imply that in the case of adiabatic fluctuations there can be no generation of fluctuation modes on super-Hubble scales. However, it is possible for entropy fluctuations to be generated on such scales, one well known example being the formation of cosmic defects during a phase transition which leads to super-Hubble structures (see e.g. \\cite{RHBrev}). Another example is the generation of axion fluctuations (see e.g. \\cite{axion}). Recently, the super-Hubble range effects of entropy fluctuations have been widely studied in the context of inflationary reheating (see e.g. \\cite{bassett,fabio,Lyth}).  The aim of this work is to investigate the process of structure formation in the non-linear regime, for inflationary models based on a single scalar field. In Section 2, we show that the presence of nonlinearities is inevitable in a wide class of inflationary models. In Section 3 we find a solution (which is exact to leading order in the gradient expansion) for nonlinear modes in the case when (but not only when) perturbations satisfy a {\\it generalized slow-rolling condition}. This condition is naturally satisfied during inflation. For our solution we can define a generalized Bardeen parameter which is constant on scales larger than the Hubble radius. As a first application (Section 4), we demonstrate that there is no amplification of long wavelength gravitational fluctuations during reheating. In Section 5 we discuss some implications of our solution for the back-reaction of perturbations on the background geometry. In Section 6, the Langevin equation for stochastic inflation \\cite{Starob} is re-derived in a way which includes the effects of gravitational (and not just matter) fluctuations, using the results of Section 3 which can be self-consistently applied in an inflating Universe with nonlinear super-Hubble modes. Finally in Section 7 we demonstrate that in the case in which the generalized slow-rolling condition is violated, the generation of vector modes is inevitable. ",
        "conclusions": "We have demonstrated that in many models of inflation, the large phase space of infrared modes leads to nonlinearities, even when the amplitude of each Fourier mode is small. We were able to find a solution of the Einstein field equations for nonlinear fluctuations which is exact to leading order in the gradient expansion, and thus will be accurate to describe the infrared modes. For this solution, we were able to define a generalized Bardeen parameter which is conserved in time.  As a first application of this formalism, we were able to show that in models with a single matter field, there can be no parametric amplification of super-Hubble cosmological fluctuations during inflationary reheating. Applied to the problem of back-reaction of infrared modes, our solution implies that the back-reaction effect is not locally identifiable as an effect due to inhomogeneities since the local Friedmann equations are satisfied. We were able to use our formalism to give a re-derivation of the equation of motion for stochastic inflation which takes into account the effects of gravitational fluctuations. Finally, we have shown that nonlinearities inevitably lead to the generation of vector perturbations if a {\\it generalized slow-rolling condition} is not satisfied.  This work has been supported in part by the US Department of Energy under Contract DE-FG02-91ER40688, Task A."
    },
    "0011/astro-ph0011518_arXiv.txt": {
        "abstract": "We present X-ray imaging spectroscopy of the extremely luminous infrared galaxy IRAS 09104+4109 ($z=0.442$) obtained with the Chandra X-ray Observatory. With the arcsec resolution of Chandra, an unresolved source at the nucleus is separated from the surrounding cluster emission. A strong iron K line at 6.4 keV on a very hard continuum is detected from the nuclear source, rendering IRAS 09104+4109 the most distant reflection-dominated X-ray source known. Combined with the BeppoSAX detection of the excess hard X-ray emission, it provides further strong support to the presence of a hidden X-ray source of quasar luminosity in this infrared galaxy. Also seen is a faint linear structure to the North, which coincides with the main radio jet. An X-ray deficit in the corresponding region suggests an interaction between the cluster medium and the jet driven by the active nucleus. ",
        "introduction": "IRAS 09104+4109 (FSC 09105+4108) is one of the extremely luminous infrared galaxies detected in the IRAS survey and has been considered to be a dust-enshrouded type 2 QSO (Kleinmann et al 1988; Hines \\& Wills 1993; Crawford \\& Vanderriest 1996; Granato, Danese \\& Franceschini 1996; Soifer et al 1996; Taniguchi et al 1997; Evans et al 1998; Hines et al 1999; Franceschini et al 2000; Tran, Cohen \\& Villar-Martin 2000). The infrared source ($L_{\\rm ir}\\simeq 6\\times 10^{12}h^{-2}$\\Ls\\footnote{\\H0\\ = 100 $h$ \\kmpspMpc\\ and \\q0\\ = 0.5 are assumed throughout this paper.}) has been identified with a cD galaxy in a rich cluster at a redshift of 0.442 (Kleinmann et al 1988). Extended X-ray emission was detected with the ROSAT HRI, and has been identified with the intracluster medium (ICM) in which a strong cooling flow is taking place (Fabian \\& Crawford 1995; Crawford \\& Vanderriest 1996; Allen \\& Fabian 1998). Although the cluster emission itself is of great interest, it makes an X-ray investigation of the active nucleus difficult for X-ray telescopes with a relatively low spatial resolution like ASCA (e.g., Fabian et al 1994) and even XMM. The first X-ray evidence for a luminous active nucleus came from an observation in the hard X-ray band where cluster emission is negligible. The detection of a hard X-ray excess with BeppoSAX, although its significance is marginal ($3\\sigma$) and subject to possible contaminating sources in the large field of view of the PDS detector (Frontera et al 1997), suggests a strongly absorbed nucleus (Franceschini et al 2000). The other way to access the active nucleus is to utilise a high spatial resolution to separate the nucleus from the diffuse emission in an appropriate energy range. With the one-arcsec resolution of the X-ray telescope (Weisskopf, O'Dell \\& van Speybroeck 1996) and the ACIS CCD detectors (Garmire et al 2000), Chandra overcomes the difficulty. In this letter, we present the spatially resolved X-ray source associated with the nucleus of IRAS 09104+4109 and its spectral properties. ",
        "conclusions": "\\subsection{Buried QSO}  A distinctive cold reflection feature is observed from the nucleus of IRAS 09104+4109 in the Chandra spectrum whilst no primary radiation is visible. Detection of the primary emission at higher energies with BeppoSAX (Franceschini et al 2000) depends, as demonstrated above, on the assumed spectrum of the source.  If the PDS-detected X-rays are entirely due to reflection, only a lower limit of the primary source luminosity can be obtained which is a minimum requirement to produce the observed luminosity through cold reflection. The albedo ($\\eta $) in the 2--10 kev band is calculated using the {\\tt pexrav} model for two inclination angles, $i=20^{\\circ}$ and 60$^{\\circ}$, for the reflecting slab subtending $2\\pi$ in solid angle. For $\\Gamma = 1.4$, $\\eta = 0.066$ ($i=20^{\\circ}$) and 0.052 ($i=60^{\\circ}$). Therefore the luminosity of the primary source is larger than $5\\times 10^{43}\\eta^{-1}\\simeq 7.6\\times 10^{44}$\\ergps\\ and $9.6\\times 10^{44}$\\ergps\\ for respective inclinations. Since the fraction of reflecting surface visible to us is likely to be less than unity, the true luminosity of the primary source should be larger than these values.  We discuss below the case in which the primary source has a power-law of $\\Gamma = 2$ and hence its transmitted radiation is detected with the PDS. The 2--10 keV luminosity of the primary source corrected for the absorption (and effects of Compton scattering, Matt et al 1999) is estimated to be $7.4\\times 10^{45}\\gamma h^{-2}$\\ergps ($\\gamma\\sim 1$ for a spherical obscuration, and it can go up to $\\sim 3$ if absorption occurs at a small cloud in the line of sight when \\nH $=3\\times 10^{24}$\\psqcm). The observed-to-intrinsic luminosity ratio in the 2--10 keV band is then $0.67\\gamma^{-1}$ per cent for IRAS 09104+4109. If the ratio of the 2--10 keV to bolometric luminosities is assumed to be $f_{\\rm HX/Bol}=0.05f_{\\rm HX/Bol,0.05}$, typical for quasars (e.g., Elvis et al 1994), the bolometric luminosity is $1.5\\times 10^{46}f_{\\rm HX/Bol,0.05}^{-1} \\gamma h^{-2}$\\ergps. This is comparable to $L_{\\rm ir}=2.3\\times 10^{46}h^{-2}$\\ergps, and consistent with a picture where the bulk of the infrared emission is powered by a buried quasar through dust reradiation. The radio source has an intermediate power between FRI and FRII (Kleinmann et al 1988; Hines \\& Wills 1993). However, without the nuclear obscuration, the wide band energy distribution of IRAS 09104+4109 would be closer to radio quiet AGNs.  The cold gas reflecting the observed X-rays is probably part of the clouds occulting the central source. An X-ray absorber with as large a column density as found in IRAS 09104+4109 is likely to be compact, probably parsec-scale in radius as for nearby Compton-thick Seyfert 2 galaxies (Matt 2000). The flat energy distribution in the infrared band indicates that hot dust at near sublimation temperature ($\\sim 1000$ K) is present and the inner radius of the dusty torus is the order of 1 pc (Granato et al 1996; Taniguchi et al 1997). The mid-infrared spectrum obtained from ISOCAM CVF {(Taniguchi et al 1997) fits well with the compact dusty torus model of Pier \\& Krolik (1992) with an edge-on view (also see Granato \\& Danese 1994; Granato et al 1996) whereas optical polarization studies imply a moderately inclined torus (Hines et al 1999; Tran et al 2000). The characteristic dust temperature is about 120 K (e.g., Kleinmann et al 1988). No IRAS detection at 100$\\mu $m and the new SCUBA limit at 850$\\mu $m (Deane \\& Trentham 2000) rule out the presence of cool dust, which is a common feature of local ULIGs. No detection of CO (Evans et al 1998) is consistent with a relatively small obscuring torus with a modest mass. A large covering factor ($\\geq 0.9$) is favoured by all the observations.   The soft X-ray emission seen in the nuclear spectrum is probably due to photoionized gas in the inner nucleus. The highly polarized biconical reflection nebula imaged by HST (Hines et al 1999) is however not a likely source, since scattering by dust rather than electrons appears to be a dominant mechanism to induce the high polarization (Tran et al 2000). The inner wall of the obscuring torus is exposed to the intense radiation from the primary source and thus expected to be highly ionized. Provided the optical depth of the ionized gas is small, the Fe-L emission bump can be very strong (e.g., Band et al 1990). A sub-parsec scale ionized disk has indeed been imaged with the VLBA in the nucleus of the nearby Compton-thick Seyfert 2 galaxy NGC1068 (Gallimore, Baum \\& O'Dea 1997), and they predicted strong soft X-ray emission lines based on a photoionization computation. In this case, the luminosity of the emission from the photoionized gas should be larger than observed ($\\sim 3\\times 10^{42}h^{-2}$\\ergps\\ in the 0.5--2 keV band), as significant absorption is likely to occur during the escape from the nuclear region.  A possible alternative is hot stars near the nucleus, inferred from a Wolf-Rayet feature seen in the optical spectrum of the nuclear region (Tran et al 2000). However, the luminosity of this component alone is one order of magnitude more luminous than the nearby Wolf-Rayet galaxies observed with ROSAT (Stevens \\& Strickland 1998), and the high metallicity implied from the strong emission-line features is also unusual (but see Buote 2000).   \\subsection{Implications for other luminous IR galaxies}  IRAS 09104+4109 is so far the only hyper-luminous infrared galaxy detected in X-rays (apart from PG1634+706 which is an unobscured quasar, see Nandra et al 1995 for the ASCA result), and the most distant reflection-dominated spectrum X-ray source known with a clear detection of the Fe K line. The small X-ray reflection fraction ($\\sim 0.7$ per cent in the 2--10 keV band; $\\sim 0.1$ per cent at the rest energy of 2 keV) estimated for IRAS 09104+4109 is still consistent with the limit obtained from the previous observations for other luminous infrared galaxies at high redshift (see Fabian et al 1997; Ogasaka et al 1997 for F15307+3252; Wilman et al 1998). Chandra and XMM are capable of detecting reflected X-ray light from F15307+3252, if the reflection fraction is similar to IRAS 09104+4109. Although a $2\\sigma $ detection of the lensed object IRAS F10214+4724 with ROSAT HRI was reported (Lawrence et al 1994), an 80 ks ASCA observation failed to detect it (Iwasawa 2000). The prospect of X-ray detection of this object depends on the lensing geometry. It should be noted that IRAS 09104+4109 is peculiar among the ULIG population being in a cluster environment, and how common this type of objects is is unclear (e.g., Deane \\& Trentham 2000).   Chandra observations of the SCUBA sources have revealed no correlation between the sub-millimetre and X-ray source populations (Fabian et al 2000b; Hornschemeier et al 2000) with two exceptions (Bautz et al 2000). If ULIGs at high redshift are powered predominantly by AGN, as suspected by Trentham, Blain \\& Goldader (1999); Trentham (2000), their nucleus is Compton-thick and the X-ray reflection fraction is indeed as low as that in IRAS 09104+4109. On the other hand, the sources with a small X-ray optical depth, that are preferably detected in X-ray, may tend to have hotter dust (Wilman, Fabian \\& Gandhi 2000) which deters submillimetre detection."
    },
    "0011/astro-ph0011032_arXiv.txt": {
        "abstract": "We present the strategy and status of a gravitational lens search for multiple imaging with angular separations between $6^{\\prime\\prime}$ and $15^{\\prime\\prime}$ within the combined JVAS and CLASS dataset of $\\sim15\\,000$ flat-spectrum radio sources. Currently all but one lens candidate have been rejected. ",
        "introduction": "The Jodrell-Bank VLA Astrometric Survey (JVAS) and the Cosmic Lens All Sky Survey (CLASS) samples together comprise $\\sim15\\,000$ radio sources in the northern sky selected to be flat-spectrum between 1.4~GHz and 5~GHz (see the contribution by Browne \\& Myers). The simple compact structures of such sources ease the task of identifying potential gravitational lensing events. So far, a search for $0.\\!\\!^{\\prime\\prime}3$ to $6^{\\prime\\prime}$ scale lensing (produced by galaxy-scale masses) has been particularly fruitful, with at least 18 gravitational lens systems found within the two samples. This poster describes a search for further lensing events on larger angular scales, corresponding to lensing by  groups and clusters of galaxies. ",
        "conclusions": ""
    },
    "0011/astro-ph0011342_arXiv.txt": {
        "abstract": "We report observation results of the super-soft X-ray sources (SSS) CAL87 and RXJ0925.7--4758 with the X-ray CCD cameras (Solid-state Imaging  Spectrometer; SIS) on-board the ASCA satellite. Because of the superior energy resolution of SIS ($\\Delta E/E \\sim 10 \\% $ at 1 keV) relative to previous instruments,  we could study detailed X-ray spectral structure of these  sources for the first time. We have applied theoretical spectral models to CAL87, and constrained the white dwarf mass and intrinsic luminosity as $0.8 - 1.2 M_\\odot $ and $4 \\times 10^{37}$--$1.2 \\times 10^{38}$ erg s$^{-1}$, respectively.  However, we have found the observed luminosity is an order of magnitude smaller than the theoretical estimate, which indicates the white dwarf is permanently blocked by the accretion disk, and we are observing a scattering emission by a fully ionized accretion disk corona (ADC) whose column density is $\\sim 1.5 \\times 10^{23}$ cm$^{-2}$. Through simulation, we have shown that the orbital eclipse can be explained by the ADC model, such that a part of the extended X-ray emission from the ADC is blocked by the companion star filling its Roche lobe.  We have found that  very high surface gravity and temperature, $\\sim 10^{10}$ cm s$^{-2}$ and $\\sim $100 eV respectively, as well as a strong absorption edge at $\\sim$ 1.02 keV, are required to explain the X-ray energy spectrum of RXJ0925.7--4758. These values are  only possible for an extremely heavy white dwarf near the Chandrasekhar limit. Although  the super soft source luminosity should be $\\sim 10^{38}$ erg s$^{-1}$ at the Chandrasekhar limit, the observed luminosity of  RXJ0925.7--4758 is nearly two orders of magnitude smaller even assuming  an extreme  distance of $\\sim$ 10 kpc. To explain the luminosity discrepancy, we propose a model that very thick matter which was previously ejected from the system,  as a form of jets, intervenes  the line of sight, and reduces the luminosity significantly due to Thomson scattering. ",
        "introduction": "Dozens of the Super Soft X-ray Sources (SSS), which  characteristically are radiating most energies in the softest X-ray energy band ($\\lesssim$ 0.5 keV), have been discovered in LMC, SMC, M31 and Milky Way by  {\\em Einstein} and ROSAT (for review, e.g., Kahabka and van den Heuvel 1997). It is believed that most SSS are white dwarf binaries with large mass accretion rates ($\\sim 1 -5 \\times 10^{-7} M_\\odot$ yr$^{-1}$), and the energy source is steady nuclear burning on the white dwarf surface (van den Heuvel et al.\\ 1992; Heise et al.\\ 1994). Observationally, our knowledge on the X-ray energy spectra of SSS has been limited by the capability of the conventional X-ray detectors.  While the energy bands of the {\\em Einstein} IPC and ROSAT PSPC (0.1 -- 2 keV) are suitable for the study of SSS, their energy resolutions ($\\Delta E/E \\gtrsim 1$ at 0.5 keV) have not allowed one to clearly identify any possible spectral structures such as emission lines or absorption edges. Having a superior energy resolution ($\\Delta E/E \\sim 10 \\%$ at 0.5 keV), the ASCA Solid State Spectrometer (SIS; Tanaka, Inoue and Holt 1994) is potentially a powerful instrument for spectroscopic study of SSS, though its detectable energy band (0.4  -- 10 keV) is slightly too high to cover the typical SSS spectra.  The two SSS, CAL87 and RXJ0925.7--475, are known to have relatively hard energy spectra, thereby we have chosen these two sources for precise spectral study with the ASCA SIS. While most SSS have characteristic blackbody temperatures of $\\lesssim$ 30 eV, those of CAL87 and RXJ0925.7--475 are $\\gtrsim$ 40 eV (Schmidtke et al.\\ 1993; Motch, Hasinger and Pietsch 1994). CAL87 was discovered during the LMC survey observations with the {\\em Einstein} IPC (Long, Helfand and Grabelsky 1981). It is an eclipsing binary with a  10.6 hour orbital period, exhibiting orbital eclipses in the optical band and accompanying shallow X-ray dips (Schmidtke et al.\\ 1993; Alcock et al.\\ 1997; Asai et al.\\ 1998). With the assumption that the non-degenerate companion is a F-star, Cowley et al.\\ (1990) speculated that the compact object is massive and may be a black hole. However, HST observations revealed two adjacent stars at 0.\\arcsec88 and 0.\\arcsec65 from CAL87 (Deutsch et al.\\ 1996), which have yet to be taken into account to determine the optical properties  of CAL87. Schandl, Meyer-Hofmeister and Meyer (1997) successfully  modeled the observed optical light curve of CAL87 assuming a 0.75 $M_\\odot$ white-dwarf and a 1.5  $M_\\odot$ companion. To account for the shallow X-ray dips, the Accretion Disk Corona (ADC) model was proposed in which the blocking material partially covers the largely extended X-ray emitting corona (Schmidtke et al.\\ 1993).  RXJ0925.7--475 was discovered in the ROSAT Galactic Plane survey project  (Motch, Hasinger and Pietsch 1994). This  unusual, heavily absorbed star exhibits most X-ray emission at energies above $\\sim$ 0.5 keV, while most other SSS have little X-rays above $\\sim$ 0.5 keV (Motch, Hasinger and Pietsch 1994; Motch 1996). The orbital period is $\\sim$ 3.8 days, and the optical counterpart indicates  strong reddening (Motch, Hasinger and Pietsch 1994; Motch 1996; Schmidtke et al.\\ 2000). A transient optical jet, similar to those observed from other super-soft sources, has been also observed (Motch 1998). The reddening, strength of the interstellar absorption lines, and ROSAT spectra are all consistent with a large hydrogen column density of $N_H \\gtrsim 10^{22}$ cm$^{-2}$, which suggests that the source is located behind the nearby ($d$=425 pc) Vela Sheet molecular cloud (Motch, Hasinger and Pietsch 1994). Optical light curve analysis suggests that mass of the compact object is in the range of 0.5 -- 1.7 $M_\\odot$ and that of the donor star is in 1 -- 2 $M_\\odot$ (Schmidtke et al.\\ 2000).  In this paper, we report precise analysis of the X-ray energy spectra of CAL87 and RXJ0925.7--4758 observed with ASCA. Owing to the superior energy resolution of the ASCA SIS, we could reveal detailed X-ray spectral features of the two sources. The ASCA CAL87 light curve exhibits shallow orbital eclipses, and its energy spectrum show deep absorption edges which are identified as \\ion{O}{7} and \\ion{O}{8} edges, as already reported by  Asai et al.\\ (1998). Preliminary spectral analysis of RXJ0925.7--4758 has been reported by Ebisawa et al.\\ (1996), and similar ASCA spectral analysis has been also carried out by Shimura (2000). As an extension of these studies, we apply theoretical  LTE (Local Thermal Equilibrium) and Non-LTE white-dwarf atmospheric spectral models to CAL87 and RXJ0925.7--4758, and constrain the white dwarf parameters.  Also, we carry out a simulation to explain the CAL87 orbital light curve in the framework of the  accretion disk corona (ADC) model, such that a part of the extended ADC X-ray emission  is blocked by a  companion star and observed as a shallow and wide  orbital eclipse. ",
        "conclusions": "We have observed the super soft sources CAL87 and RXJ0925.7--4758  with ASCA SIS, and carried out precise spectral analysis. Thanks to the superior energy resolution of SIS ($\\Delta E/E \\sim 10 \\% $ at 1 keV) compared to previous instruments,  we could study detailed X-ray spectral structure of these  sources for the first time. Important results are summarized in the following:  (1) We have applied theoretical spectral models to CAL87, and constrained the white dwarf mass and intrinsic luminosity as $0.8 - 1.2 M_\\odot $ and $4 \\times 10^{37}$--$1.2 \\times 10^{38}$ erg s$^{-1}$, respectively.  We have found the observed luminosity is an order of magnitude smaller than the intrinsic luminosity, which indicates the white dwarf is permanently blocked by the accretion disk and not directly seen. This strongly suggests that  we are observing a scattering emission by a fully ionized accretion disk corona (ADC) whose column density is $\\sim 1.5 \\times 10^{23}$ cm$^{-2}$.  (2) Through simulation, we have shown that the orbital eclipse can be explained by the ADC model, such that a part of the extended X-ray emission from the ADC, whose size is expected to be  $\\sim$ 75 \\% of the accretion disk radius, is blocked by the companion star filling its Roche lobe. The orbital inclination angle is  $\\sim 73^\\circ$ to explain the observed eclipse profile.  (3) In order to explain the RXJ0925.7--4758 spectrum, very high surface gravity and temperature, $\\sim 10^{10}$ cm s$^{-2}$ and $\\sim $100 eV respectively,  are required. These values are  only possible for an extremely heavy white dwarf near the Chandrasekhar limit.  Therefore, RXJ0925.7--4758 may be an immediate progenitor of Type Ia supernovae, which has been  long-sought (e.g., Livio et al.\\ 1996). We have found the energy spectrum has a $\\sim$ 1.02 keV edge, which may be a composite of mildly ionized L- or K-edges of heavy elements.  (4) Although  super-soft sources should have the  luminosity $\\sim 10^{38}$ erg s$^{-1}$ at the Chandrasekhar limit, the observed luminosity from RXJ0925.7--4758 is smaller by more than two orders of magnitude. To explain the luminosity discrepancy, we propose a model that very thick matter (as much as  $\\sim 10^{25}$ cm$^{-2}$) which was previously ejected from the system as jets intervenes  the line of sight, and reduces the luminosity significantly due to Thomson scattering.  We have shown that the matter is almost fully ionized that it is transparent for soft X-ray photons. Some  radial portion of the ejected matter, where ionization state is locally low and heavy elements are not fully ionized, may be origin of the significant  low energy absorption and the  $\\sim$ 1.02 keV edge.  We wish to thank the anonymous referee for his/her carefully reading the manuscript and useful comments."
    },
    "0011/astro-ph0011497_arXiv.txt": {
        "abstract": "We have used an updated version of the empirically and semi-empirically calibrated \\basel\\ library of synthetic stellar spectra of Lejeune \\etal\\ (1997, 1998) and Westera \\etal\\ (1999) to calculate synthetic photometry in the \\ubvetc,  HST-WFPC2, Geneva, and  Washington systems for the entire   set of non-rotating  Geneva stellar  evolution models covering  masses from 0.4--0.8 to   120--150 \\msun\\ and  metallicities $Z$=0.0004 (1/50 \\zsun) to 0.1 (5 \\zsun). The results are provided  in a database\\footnote{The full database will be available in electronic form at the CDS {\\tt http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/(vol)/(page)} and at {\\tt http://webast.ast.obs-mip.fr/stellar/}.} which includes all individual stellar tracks and   the corresponding isochrones covering ages   from 10$^3$ yr to 16--20 Gyr in time steps of $\\Delta \\log t=$ 0.05 dex.  The database also includes a new grid of stellar tracks of very metal-poor stars ($Z$=0.0004)  from 0.8 --  150 \\msun\\ calculated with the Geneva stellar evolution code. ",
        "introduction": "Studies   of individual  stars  and resolved   stellar populations are fundamental  for a wide  variety  of astrophysical subjects.  In  most cases photometric observations, obtained in various systems, represent the main observable. To translate these observables into fundamental stellar parameters (mainly effective  temperature, luminosity) one first has to rely on stellar atmosphere models. In a second step the observations can then be compared to stellar evolution models and interpreted in physical terms such as mass, age, composition etc. Only in rare cases (very  low mass stars,  Baraffe \\etal\\ 1995; O stars, Schaerer \\etal\\ 1996) stellar evolution and atmosphere models are coupled and predict thus directly observables such as colors, magnitudes etc.  Extensive  grids    of stellar tracks  covering    the  most important evolutionary phases   and   a large  metallicity   range  have  become available  in the nineties  (see e.g.\\ compilation in Leitherer \\etal\\ 1996).  In  addition, recent efforts have   been undertaken to provide accurate synthetic  colours  from grids of  atmosphere models covering the bulk of the parameter space occupied by observed stars of spectral types from O  to M and  all luminosity  classes (Lejeune \\etal\\  1997, 1998, Bessell \\etal\\ 1998). With   the availability  of   such data it    now becomes  feasible to systematically convert stellar    tracks  of all stellar  masses   and derived isochrones  to a variety  of photometric  systems.  To provide such a tool is the main goal of our database.  Existing grids, partly fulfilling this aim, include the library of Padova isochrones with UBVRIJHK photometry for ages $\\sim$ 4 Myr -- 20 Gyr and metallicities $Z$ between 0.0004 and 0.05 (Bertelli \\etal\\ 1994, Girardi \\etal\\ 1996), the recent $\\alpha$-element enhanced tracks and isochrones of Salasnich \\etal\\ (2000) with UBVRIJHK and HST-WFPC2 photometry, and various other calculations covering smaller fractions of the parameter space.  For the present work we rely on an updated version of the hybrid library of synthetic stellar spectra compiled by Lejeune \\etal\\ (1997, 1998), which is corrected to match empirical colour--temperature relations at solar metallicity and semi-empirically corrected at other metallicities. These atmosphere models are used to derive synthetic photometry of the \\ubvetc, WFPC2, Geneva, and Washington systems, using proper reference spectra for the zero points and up-to-date filter curves. This procedure is applied to essentially the full set of Geneva stellar evolution models to yield the photometric data for these systems for {\\em 1)} individual stellar tracks in the mass range from 0.4--0.8 to 120--150 \\msun\\ and metallicities $Z$=0.0004 (1/50 \\zsun) to 0.1 (5 \\zsun), and {\\em 2)} corresponding isochrones covering ages from $\\sim$ 0 to 16--20 Gyr.  The input tracks, atmosphere models,  and the synthetic photometry are described in Sect.\\ 2.  The  resulting database products  (tracks  and isochrones) are presented in  Sect.\\ 3. Concluding  remarks on the use of the  various database sets  are given in  Section  4.  The Appendix includes the relevant data for a previously unpublished set of stellar tracks at $Z$=0.0004. ",
        "conclusions": ""
    },
    "0011/astro-ph0011174_arXiv.txt": {
        "abstract": "s{The probability that high-redshift sources are gravitationally-lensed with large image separations (\\ie, greater than can be produced by galactic deflectors) is determined by the cosmological population of group- and cluster-sized halos. Thus the observed frequency of wide-separation lensed quasars can be used to constrain not only the halo distribution, but also a number of cosmological parameters. A calculation of the optical depth due to collapsed, isothermal halos is a useful guide to the lens statistics, and illustrates that the number of wide-separation lenses is a sensitive probe of the mean density of the universe and the present day density variance whilst being nearly independent of the cosmological constant.} ",
        "introduction": "\\label{section:intro}  Gravitational lensing is a direct probe of the geometry of the universe, and so observations of lensing place constraints on both the distribution of mass and the underlying cosmological model. A large variety of lensing experiments are possible: weak shear surveys; searches for microlensing by compact objects; observations of giant arcs in clusters; measurement of Shapiro delay in binary pulsars; surveys for multiply-imaged qusasrs; and so on. The focus here is on the statistics of wide-separation lensed quasars (which are produced by group- and cluster-mass objects). Using a simple model of the halo population (\\sect{deflectors}), the optical depth to multiple imaging and the expected distribution of image separations can be calculated (\\sect{calc}). The results obtained are independent of the source population, which then implies that a more sophisticated analysis is required to constrain model parameters from the observed frequency wide-separation lenses~\\cite{ko95}$^{,}$~\\cite{mo00}. ",
        "conclusions": "\\label{section:concs}  The statistics of wide-separation lenses are a useful cosmological probe. The lensing probability is sensitive to the geometry of the universe, and the population of collapsed halos at intermediate redshifts whilst being independent of the complexities of galaxy formation. As large-scale structure is dependent on the linear growth of perturbations (and also because the mass in a co-moving sphere of a given size is proportional to the density parameter) the likelihood of lensing is most sensitive to $\\Omega_{\\rm m_0}$ and $\\Delta_8$. The Large Bright Quasar Survey~\\cite{he95} is devoid of wide-separation lenses~\\cite{he98}$^{,}$~\\cite{mo99} which implies that $\\Omega_{\\rm m_0} \\la 0.3$ (assuming $\\Delta_8 \\ga 0.4$) or that $\\Delta_8 \\la 0.6$ (assuming $\\Omega_{\\rm m_0} \\la 0.1$) with 99 per cent confidence~\\cite{mo00}. In the future such analyses could be extended to the 2 degree Field quasar survey~\\cite{bo99} (with $\\sim 3 \\times 10^4$ quasars) and the Sloan Digital Sky Survey~\\cite{sz98} (which will contain $\\sim 10^5$ sources). These samples will be large enough to constrain $\\Omega_{\\rm m_0}$ and $\\Delta_8$ to within several per cent from the number of wide-separation lenses alone."
    },
    "0011/hep-ph0011261_arXiv.txt": {
        "abstract": "Hybrid inflation is a natural scenario in the absence of supersymmetry. In the context of supergravity, however, it has to face the naturalness problems of the initial conditions and of the adequate suppression of the inflaton mass. Both can be successfully addressed in a class of models involving K\\\"ahler potentials associated with products of $SU(1,1)/U(1)$ K\\\"ahler manifolds and ``decoupled'' fields acquiring large vacuum expectation values through D-terms. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011424_arXiv.txt": {
        "abstract": "Stars are most likely to merge or collide in regions of the highest stellar densities, and our own Galactic Center contains many stars packed into a relatively small volume -- even the ambient stellar number density in the central 50~{\\it pc} is quite high, $\\sim10^3~stars~pc^{-3}$. More striking, the three compact young clusters in this region have central densities as high as $10^6~stars~pc^{-3}$. We discuss these extreme environments and the possibility that stellar mergers and collisions have recently occured there. In particular, we predict that at least one massive star in the Arches cluster has already experienced a stellar merger in its short lifetime. Further, the Pistol Star, in the nearby Quintuplet cluster, might owe its apparent relative youth to a rejuvinating stellar merger. Finally, the apparently young stars in the central arcsecond could be products of either collisions, inducing atmospheric stripping, or mergers. ",
        "introduction": "The Galactic Center (GC) is a very dense stellar environment and was recognized early on as a region which might contain products of stellar mergers and collisions. Soon after the identification of the dense stellar cluster in the GC (Becklin et al.\\ 1978), Lacy, Townes, \\& Hollenbach (1982) suggested that material in the ionized gas in the GC mini-spiral was produced by star-star collisions, but it is now thought that the mini-spiral traces gas pulled in from a wayward molecular cloud, or the nearby circumnuclear ring. Nonetheless, stellar collisions and mergers may still play an important role in describing phenomena in the central parsec. Indeed, a new favorite target for this possibility is the tight cluster of stars in the central arcsecond (Bailey \\& Davies 1998; Alexander 1999), as its high velocity stars appear (Genzel et al.\\ 1997; Ghez et al.\\ 1998) to be very young (Eckart et al.\\ 1999; Figer et al.\\ 2000).  There are, at least, two other interesting regions, as regards stellar collisions, just 30~{\\it pc}, in projection, from the center: the Arches cluster (Nagata et al.\\ 1995; Cotera et al.\\ 1996; Serabyn, Shupe, \\& Figer 1997; Figer et al.\\ 1999a), and the Quintuplet cluster (Okuda et al.\\ 1990; Nagata et al.\\ 1990; Glass, Moneti, \\& Moorwood 1990; Figer, McLean, \\& Morris 1999b). They are nearly identical to each other, and to the young cluster in the central parsec, except for age -- the Arches cluster is $\\sim$2~Myr old, half the age of the other two clusters.  This paper is an introduction to the GC, and objects therein, as it pertains to studies of stellar collisions and mergers. We start by introducing the environment of the GC and finish with 3 case studies of potential mergers or collisions. ",
        "conclusions": "In conclusion, we find that the three young stellar clusters in the Galactic Center are prime proving grounds for stellar merger and collision theories. In particular, these clusters are the densest and most massive young clusters in the Galaxy. They have individual members which might be products of stellar mergers, i.e.\\ the Pistol Star, and they may contain products of near-misses, or collision, i.e.\\ in the Central cluster."
    },
    "0011/astro-ph0011562_arXiv.txt": {
        "abstract": "{We present new dynamical parameters of the AR~Pav binary system. Our observations consist of a series of high resolution optical/NIR spectra from which we derive the radial velocity curve of the red giant as well as its rotation velocity. Assuming co-rotation, we determine the stellar radius (130\\,\\Rsolar) of the red giant. Based on this we derive the red giant's luminosity and mass (2.0\\,\\Msolar) as well as the distance of the system (4.9\\,kpc). The binary mass function finally yields the companion's mass (0.75\\,\\Msolar) and the binary separation (1.95\\,AU). We find that the red giant does not fill its Roche lobe. We review the radial velocity data of Thackeray and Hutchings (1974), and compare it with our red giant's orbit. We find that their RV curves of the blue absorption system and the permitted emission lines are in anti-phase with the red giant, and that the forbidden emission lines are shifted by a quarter of a period. The blue absorptions and the permitted emission lines are associated with the hot companion but not in a straightforward way. The blue absorption system only tracks the hot component's orbital motion whilst it is in front of the red giant, whereas at other phases line blanketing by interbinary material leads to perturbations. We finally present UV light curves based on IUE archive spectra. We clearly detect eclipses in the continuum at all wavelengths. The eclipse light curves are unusual in that they show a slow and gradual decline prior to eclipse which is followed by a sharp increase after eclipse. ",
        "introduction": "AR Pavonis was the first symbiotic star to be recognized as a binary system. Mayall (1937) discovered it as an eclipsing P Cygni type star with a period of 605 days. She pointed out, how remarkable it was to observe an eclipsing star with such a long period. She also warned that its variability was unusual for a P Cygni star. Indeed, the association with P Cygni was only due to the presence of Balmer lines in emission with strong absorption edges toward the violet. As MWC 600 it entered Merrill and Burwell's (1943) catalogue of Be and Ae stars as type Beq. Although Sahade (1949) extended the list of observed absorption and emission lines, it was Thackeray (1959) who made a first attempt at modelling, breaking away from the idea of a peculiar P Cygni variable. Finally, Thackeray \\&\\ Hutchings (1974) claimed AR Pav for the symbiotic family. Based on their extended set of spectroscopic observations, covering several eclipses, they proposed a binary model. The strength of the TiO bands suggested an M3 III as secondary. In their opinion the eclipse was not just due to the extinction of stellar radiation, but to the ``eclipse of a region whose excitation and density increases towards a central source of light''. They distinguish different radial velocity behaviours in permitted emission, forbidden emission, and absorption lines. Permitted emission lines, e.g. \\ HeII $\\lambda4686$, give the clearest eclipse pattern. They associate it with nebular emission centered on what to them is the hot primary, and they derive a nearly circular orbit and a mass function of $m\\tief f\\hoch {hot}$=0.14. Based on considerations about the eclipse length and the probable mass of the secondary, they estimate a mass ratio of $q\\sim$~2 and suggest a hot ($T\\approx30\\,000$\\,K) star with a mass of $\\approx2$\\,\\Msolar, or more precisely $M\\tief{hot}= 2.5$\\,\\Msolar, $M\\tief{cool}= 1.2$\\,\\Msolar.  The same mass function was adopted by Skopal et al.~(2000) when they analysed the 1889-1998 light curve to arrive at $M\\tief{hot}= 4.5$\\,\\Msolar, $M\\tief{cool}= 1.8$\\,\\Msolar. Thackeray \\&\\ Hutchings (1974) do find less clear velocity patterns in the forbidden emission or the absorption lines in the blue continuum. They associate them with gas streams (see Fig.\\,5 in Thackeray \\& Hutchings), and do not derive a velocity curve for the cool giant. In this paper we question the assumption of Thackeray \\& Hutchings (1974) that the allowed emission lines accurately trace the orbit of the hot star.  We have observed a set of absorption lines from the red giant (Fig.~\\ref{spec.ima}) over 3.5 binary periods. From this we derive an orbital solution, including a revised mass function. Our observations also permit us to extract the rotation velocity of the red giant, and with the assumption of corotation we derive the absolute radius of the red giant, its luminosity, and the distance to AR Pav.  The analysis of Thackeray \\& Hutchings (1974) as well as our analysis are based on the assumption that the occurence of an eclipse implies an approximate coincidence of the line of sight with the binary orbital plane of AR Pav. However, this has been contested by Hutchings et al. (1983). They claim that IUE (International Ultraviolet Explorer) observations prove that the hot star is not eclipsed. We re-anaylize the IUE data and establish that eclipses are clearly present.  Our orbital analysis of the red giant spectrum follows the same pattern as was already applied to the symbiotic systems SY~Mus (Schmutz et al.\\ 1994, Paper~I), RW~Hya (Schild et al.\\ 1996, Paper~II), CD--43.14304 (Schmid et al.\\ 1998, Paper~III), BX~Mon (Dumm et al.\\ 1998, Paper~IV), and FG~Ser (M\\\"urset et al. 2000, Paper~V).   \\begin{figure} \\hbox{\\psfig{figure=MS10069f1.ps,height=6.0cm,width=8.5cm,clip=}} \\hskip-0.5cm \\caption[]{Example of a normalised red spectrum of AR~Pav that served for cross correlation with an RV standard star.} \\label{spec.ima} \\end{figure}   \\begin{table} \\caption[]{Measured radial velocities and log of observations for AR~Pav. The phase $\\phi$ is calculated from the orbit solution given in Table~\\ref{bahnparam}. RV is the measured heliocentric radial velocity of the cool giant. O--C are the residuals in [km/s]. The Julian date  is relative to 2\\,400\\,000.} \\label{obslog} \\begin{center} \\begin{tabular}{rrrrcccr} \\hline\\noalign{\\smallskip} \\multicolumn{3}{c}{Date} & JD~~ & $\\lambda\\tief c$ & $\\phi$ & RV & O--C\\\\ d & m &\tyr&& [\\AA]     & &[km/s] &\\\\ \\noalign{\\smallskip}\\hline\\hline\\noalign{\\smallskip} 12& 9&91 &  48511.6 & 7005 & 0.62 & $ -75.5 \\pm 1.0$ &--0.7\\\\ 24&10&   &  48553.8 & 6565 & 0.69 & $ -77.4 \\pm 1.0$ &--0.2\\\\ 15& 5&92 &  48757.7 & 6565 & 0.02 & $ -68.3 \\pm 1.0$ &--1.3\\\\ 24& 5&93 &  49131.9 & 6565 & 0.64 & $ -77.9 \\pm 1.0$ &--2.0\\\\ 29&10&   &  49289.5 & 6565 & 0.90 & $ -74.7 \\pm 1.0$ &--0.8\\\\ 31&10&   &  49291.5 & 6825 & 0.91 & $ -73.6 \\pm 1.0$ &0.1\\\\ 1&11&   &  49292.5 & 5880 & 0.91 & $ -75.0 \\pm 1.0$ &--1.4\\\\ 17& 7&94 &  49550.6 & 6565 & 0.34 & $ -61.4 \\pm 1.0$ &--1.3\\\\ 12& 8&   &  49576.6 & 7505 & 0.38 & $ -60.6 \\pm 1.5$ &1.1\\\\ 13& 8&   &  49577.7 & 8170 & 0.38 & $ -62.7 \\pm 1.5$ &--0.9\\\\ 1&11&   &  49657.6 & 7005 & 0.51 & $ -70.4 \\pm 1.0$ &--1.2\\\\ 24& 5&95 &  49861.8 & 6565 & 0.85 & $ -77.4 \\pm 0.9$ &--1.2\\\\ 29& 5&   &  49866.8 & 7453 & 0.86 & $ -74.7 \\pm 0.5$ &1.2\\\\ 16&11&   &  50037.5 & 7005 & 0.14 & $ -61.8 \\pm 1.0$ &--0.8\\\\ 17&11&   &  50038.5 & 6825 & 0.14 & $ -60.7 \\pm 0.7$ &0.2\\\\ 31& 3&96 &  50173.9 & 6565 & 0.37 & $ -58.2 \\pm 1.5$ &3.0\\\\ 2& 7&   &  50266.7 & 5007 & 0.52 & $ -70.0 \\pm 0.6$ &--0.4\\\\ 22& 7&   &  50286.6 & 5007 & 0.55 & $ -71.6 \\pm 0.7$ &--0.1\\\\ 7& 8&   &  50302.7 & 5007 & 0.58 & $ -71.0 \\pm 0.6$ &2.0\\\\ 7& 8&   &  50302.8 & 4363 & 0.58 & $ -73.6 \\pm 1.4$ &--0.6\\\\ 8& 8&   &  50303.7 & 7005 & 0.58 & $ -73.6 \\pm 0.6$ &--0.5\\\\ 20& 7&97 &  50649.8 & 5007 & 0.15 & $ -57.8 \\pm 1.5$ &2.7\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "}  We present in this paper the first direct radial velocity measurements of the red giant in the AR~Pav system. The red giant's orbit is close to circular and yields a mass function of 0.055\\,\\Msolar. We also measure the red giant's rotation velocity and, assuming co-rotation, deduce its radius from which the luminosity and distance immediately follow. Based on RGB/AGB stellar evolutionary tracks we obtain the red giant's mass, and from the above mass function the mass of the hot companion. It turns out that the AR Pav stellar components have masses similar to other symbiotic systems. In particular the mass of the companion is close to the canonical value of single white dwarfs. The red giant does not fill its Roche lobe.  The velocity amplitude of the hot component can be obtained from historically recorded radial velocities of absorption lines in the blue continuum, taken during an outburst. Combining that result with our new radial velocities of the red giant yields directly the masses of the binary components. These masses are in very close agreement with those that we derive with our co-rotation method developed and applied in Papers I to V, as well as in this work (Section 3 and 4). We therefore have now the first independent confirmation that for symbiotics, this method leads to valid results.  Our masses for the AR~Pav binary components differ strongly from earlier values found in the literature. Thackeray and Hutchings (1974) favoured a value of approximately 2.5\\,\\Msolar\\ for the hot companion and 1.2\\,\\Msolar\\ for the red giant. They assumed that the radial velocity curve of the permitted emission lines reflected the motion of the hot companion. Although this emission is certainly associated with that object, it is not clear how exactly they are linked. The interpretation of emission lines involves a number of complicating factors like self-absorption which may be phase dependent, and unknown nebular motions which may overlay the binary orbit. In addition, different nebular emission regions might contribute with variable fractions to the total intensity. We conclude that in symbiotic systems, emission lines are rather poor tracers of stellar orbital motions. It is indeed quite instructive to see that only direct observations of the stellar absorbtion features of the M star enable a proper interpretion of emission line radial velocities.  An even higher mass for the hot companion was suggested by Skopal et al. (2000), working with the same mass function as Thackeray \\& Hutchings (1974). They assumed that the system undergoes Roche lobe overflow. However, we have shown that the mass function derived from emission lines is most likely seriously wrong, and that the Roche lobe assumption is very questionable.     \\begin{table} \\caption[]{Roche lobe filling factors for symbiotic systems (see text).} \\label{roloff} \\begin{flushleft} \\begin{tabular}{lrcrcl} \\hline\\noalign{\\smallskip} \\ts\\ts Object & Period & Separation & $R\\tief c$~     & RL filling& Ref  \\\\ & [d]~~   &  [AU]      & [\\Rsolar]& factor \\\\ \\hline\\noalign{\\smallskip} \\ts\\ts AR Pav   & 604.5     & 1.95  & 130 & 0.52 & 8) \\\\ \\ts\\ts FG Ser   & 650~~     & 1.95  & 105 & 0.41 & 7) \\\\ \\ts\\ts SY Mus   & 624.5     &  1.7  & 86  & 0.38 & 4) \\\\ \\ts\\ts EG And   & 482.2     &  1.5  & 75  & 0.37 & 3) \\\\ \\ts\\ts RW Hya   & 370.4     &  1.3  & 60  & 0.34 & 5) \\\\ \\ts\\ts AX Per   & 680.8     &  1.7  & 110 & 0.30 & 1) \\\\ \\ts\\ts AG Peg   & 816.5     &  2.5  & 85  & 0.25 & 2) \\\\ \\ts\\ts BX Mon   & 1401~~    & 2.0-5.9 & 160 & 0.68-0.18  & 6) \\\\ \\hline\\noalign{\\smallskip} \\end{tabular} \\smallskip 1) Mikolajewska \\& Kenyon (1992)\\\\ 2) Kenyon et al.\\ (1993)\\\\ 3) Vogel et al.\\ (1992)\\\\ 4) Paper I\\\\ 5) Paper II\\\\ 6) Paper IV\\\\ 7) Paper V\\\\ 8) this paper \\vskip-0.4cm \\end{flushleft} \\end{table}   One of the fundamental questions in relation to symbiotic binaries is the process of mass transfer. Is it predominantly through wind accretion or do the red giants fill their Roche lobes? We now have a small, but already significant sample of symbiotic systems with accurate orbits, they allow to draw some preliminary conclusions. In Tab.~\\ref{roloff} we have summarised what we consider to be the presently best available values. The Roche lobe filling factor is the ratio of the red giant's radius to the distance to the inner Lagrange point. We note that none of the symbiotic red giants fills its Roche lobe. A similar result was already obtained by M\\\"urset \\& Schmid (1999), who analysed a larger sample of symbiotic stars with, however, less accurately known red giant radii and binary configurations. We therefore conclude that mass transfer occurs predominantly through wind accretion.  Related to this accretion process is the question whether the accretor is surrounded by an accretion disk. To model the wind accretion onto the white dwarf it is necessary to know the physical conditions in its neighbourhood. Among the crucial parameters are the white dwarf radiation field as well as its magnetic field. A strong enough white dwarf radiation field could prevent accretion, a magnetic field could lead to accretion downflows over the polar regions. For AR~Pav, unfortunately nothing is known about the white dwarf luminositiy and on its magnetic field. Wind accretion models which neglect both effects such as those of Mastrodemos \\& Morris~1998 and Mastrodemos \\& Morris~1999 suggest the formation of accretion disks for a wide range of geometric configurations and red giant wind parameters. Some of their models are similar to those of AR~Pav. With the observed radial velocity behaviour of the permitted emission lines, it is likely that at least part of these emission lines originate in the vicinity of the white dwarf. Applying eclipse techniques to available and new high resolution emission line profiles with good orbital coverage, to search for spectral traces of an accretion disk and to study the circumstellar matter distribution would certainly be an interesting continuation to this work."
    },
    "0011/astro-ph0011048_arXiv.txt": {
        "abstract": "Properties of nearby dwarf galaxies are briefly discussed. Dwarf galaxies vary widely in their star formation histories, the ages of their subpopulations, and in their enrichment history. Furthermore, many dwarf galaxies show evidence for spatial variations in their star formation history; often in the form of very extended old populations and radial gradients in age and metallicity.  Determining factors in dwarf galaxy evolution appear to be both galaxy mass and environment. We may be observing continuous evolution from low-mass dwarf irregulars via transition types to dwarf spheroidals, whereas other evolutionary transitions seem less likely. ",
        "introduction": "Dwarf galaxies are the most frequent type of galaxies in nearby galaxy groups and clusters, and probably in the Universe.  They are the primary building blocks of more massive galaxies in hierarchical clustering scenarios and are believed to have been even more numerous at earlier times.  The best studied dwarf galaxies are nearby dwarfs where individual stars can be resolved and evolutionary histories can be derived in great detail.  Hence the most detailed information is available for our immediate neighbors, the dwarf satellites of the Milky Way.  A growing body of data is becoming available also for more distant dwarf galaxies in the Local Group and beyond owing largely to observations with the {\\em Hubble Space Telescope} (HST).  Apart from deep observations of selected galaxies our ongoing HST snapshot survey of nearby galaxies is rapidly increasing the number of dwarf galaxies within 5 Mpc for which resolved upper-red-giant-branch photometry is available. Furthermore, observations with large 8m -- 10m ground-based telescopes such as Keck, Gemini, Subaru, and the Very Large Telescope (VLT) are playing an increasingly important role.  For the purpose of this review we will consider all galaxies with absolute $V$ magnitudes $M_V > -18$ mag and with diameters of a few kpc or less as dwarf galaxies. Dwarf galaxies are important not only as future constituents of larger galaxies, but also in their own right.  They span a wide range of masses, luminosities, mean metallicities, gas content, kinematic properties, and mean ages.  Dwarf galaxies are found in different environments ranging from voids to loose groups and dense galaxy clusters, and from relative isolation within groups to close proximity of massive galaxies.  Dwarf galaxies allow us therefore to study the impact of external environmental effects and of internal properties such as gas content and galaxy mass on galaxy evolution.  Knowledge of their internal kinematics may help to understand the nature of dark matter, and their global kinematics make them valuable dynamical probes of the dark halos of massive galaxies and clusters.  The existing, detailed studies have revealed a surprising diversity in the star formation histories of dwarf galaxies.  Each galaxy shows evidence for a complex evolution that is clearly distinct from the single-age, single-metallicity history characterizing a typical globular cluster. Moreover, even within the same morphological subclass no two dwarf galaxies are alike and differ considerably in their enrichment histories and/or time and length of their star formation episodes.  However, common global properties are beginning to emerge, and this review will concentrate on identifying them rather than on describing the differences.  For recent reviews that describe the properties of Local Group dwarf galaxies in great detail we refer to van den Bergh (1999; 2000). ",
        "conclusions": ""
    },
    "0011/astro-ph0011576_arXiv.txt": {
        "abstract": "In this paper and a companion paper we present the \\dm model, a Monte Carlo radiative transfer code, self-consistently including dust heating and emission, and accounting for the effects of the transient heating of small grains.  The code is completely general; the density structure of the dust, the number and type of heating sources, and their geometric configurations can be specified arbitrarily within the model space.  Source photons are tracked through the scattering and absorbing medium using Monte Carlo techniques and the effects of multiple scattering are included.  The dust scattering, absorbing, and emitting properties are calculated from realistic dust models derived by fitting observed extinction curves in Local Group galaxies including the Magellanic Clouds and the Milky Way.  The dust temperature and the emitted dust spectrum are calculated self consistently from the absorbed energy including the effects of temperature fluctuations in small grains.  Dust self-absorption is also accounted for, allowing the treatment of high optical depths, by treating photons emitted by the dust as an additional heating source and adopting an iterative radiative transfer scheme. As an illustrative case, we apply the \\dm radiative transfer code to starburst galaxies wherein the heating sources are derived from stellar evolutionary synthesis models. Within the context of the starburst model, we examine the dependence of the ultraviolet to far-infrared spectral energy distribution, dust temperatures, and dust masses predicted by \\dm on variations of the input parameters. ",
        "introduction": "\\label{sec:intro} Over the past two decades, studies of galaxies have become increasingly more quantitative as powerful new instruments sensitive from the far-ultraviolet (UV) to the infrared (IR) have become available.  With these observations, it has become clear that the presence of dust has a significant effect on the observed properties of galaxies; the observed spectral energy distribution (SED) is a complex convolution of the intrinsic SED of the stellar populations with the physical properties of the absorbing and scattering medium (dust), including its composition as well as its geometric relation to the stellar sources.  In addition to complicating the interpretation of observations of galaxies, dust is also an essential component in determining and modifying the physical conditions in the interstellar medium (ISM) of galaxies, regulating star formation, and participating in the chemical evolution of the galaxy.  The effects of dust are particularly pronounced in galaxies undergoing active star formation, e.g., starburst galaxies.  A substantial fraction of nearby galaxies ($\\sim$30\\%; Salzer et al. 1995) is made up of active star forming galaxies and nearly all galaxies at high redshift display characteristics typical of local star forming galaxies \\citep{hetal98}.  Therefore, the ability to quantify the effects of dust in interpreting galaxy observations has implications not only for the study of starburst galaxies themselves, but also the formation and evolution of galaxies over the age of the universe.  Quantifying the effects of dust in any astrophysical system is complicated by the geometric relationship between the illumination source(s) and the dust, uncertainty in the dust composition, and the structure of the scattering and absorbing medium.  These effects are especially pronounced in galaxies, where a typical observing aperture may include multiple, complex regions comprised of stars, gas, and dust mixed together in complicated geometries and widely varying environments.  Differences in stellar populations, metallicities, dust properties, and geometry can produce similar effects on observed properties of galaxies.  Disentangling the effects of these various intrinsic galactic properties requires the use of realistic models of the transfer of radiation including both stars and dust.  Historically, the treatment of dust in radiative transfer simulations has been necessarily simplistic.  In many cases, the dust distribution is assumed to take the form of a foreground screen in analogy with stellar extinction studies.  However, the extension of this simple geometry to more complicated systems like galaxies can lead to severely erroneous results regarding the amount of dust present and its effects on the observed SED (e.g., Witt, Thronson \\& Capuano 1992 and references therein).  Observationally, the interstellar medium in the Milky Way and external galaxies possesses structure over a range of spatial scales, characterized by variations in density over several orders of magnitude.  Finally, the physical characteristics of the dust grains determine how the grains absorb, scatter, and re-emit stellar photons.  Many models of the transfer of radiation through dusty media in galaxies have appeared in the literature and although they have included all of these factors to some degree, owing to the complexity of the problem, none treat them all simultaneously.  Some are restricted to geometries that exhibit global symmetries \\citep{err95,mh98,sgbd98,vd99,tav99}, while others assume constant or continuously varying, homogeneous dust distributions \\citep{wtc92,err95,mh98,sgbd98,tav99} or do not fully treat the re-emission from grains in the infrared \\citep{wtc92,err95,tav99,whs99}.  One of the seminal works in establishing the importance of geometry in radiative transfer models of galaxies is that of \\citet{wtc92}.  Subsequent work established the importance of local structure, i.e.  clumpiness, in modeling the transport of radiation in dusty media \\citep{wg96,wg00}.  The models developed by these authors employ Monte Carlo techniques to solve the radiative transfer equations through inhomogeneous dusty media with arbitrary geometries.  The models are quite general, including the effects of multiple scattering and non--isotropic scattering functions. The use of Monte Carlo techniques in the radiative transfer problem allows the treatment of arbitrary 3-dimensional geometries with no symmetries and can easily include a non--homogeneous, clumpy structure for the scattering and absorbing medium. However, the models do not include the effects of dust heating and re-emission. In this paper and a companion paper \\citep{gmwc00}, we present the \\dm model which incorporates the strengths of the previous models and extends them to include dust heating and re--emission in the IR.  The importance of the IR in understanding galaxies can be readily seen by considering the absorption, scattering, and emission properties of dust.  The absorption, scattering, and re-emission of photons by dust grains occur in different wavelength regimes. Dust is very efficient at absorbing and scattering UV/optical photons.  The energy absorbed by the dust is thermalized and re--emitted at IR wavelengths.  As a result, a large fraction (approaching 100\\% for heavily enshrouded regions) of a galaxy's UV/Optical energy may be reprocessed and re--emitted by the dust at IR wavelengths. Thus studies of galaxies that consider only the UV/optical wavelengths can neglect a large fraction of the galaxies' energy budgets.  A successful model of the transfer and emission of radiation in a galaxy must consistently reproduce the observed SED of the galaxy from the UV to the far infrared (FIR) simultaneously.  Reprocessing of UV/optical photons into IR photons occurs through two basic mechanisms depending both on the radiation field they are exposed to and the radius, $a$, of the dust grain.  Large dust grains reach thermal equilibrium and emit as modified blackbodies with an equilibrium temperature, $T_{eq}$. However, small grains (and also large molecules e.g., polycyclic aromatic hydrocarbons or PAH) have small heat capacities and the absorption of even a single UV/optical photon can substantially heat the grain.  These small grains will not reach an equilibrium temperature but will instead undergo temperature fluctuations that lead to grain emission at temperatures well in excess of $T_{eq}$.  The inclusion of small, thermally fluctuating grains in dust models is necessary to explain the excess of near and mid-IR emission observed in a variety of systems, including galaxies (e.g., Sellgren 1984; Helou 1986; Boulanger \\& P\\'{e}rault 1988; Sauvage, Thuan \\& Vigroux 1990). In addition, the observation of prominent emission lines widely ascribed to PAH molecules in the spectra of galaxies (e.g., Acosta-Pulido et al.\\ 1996), indicates that a realistic dust model should include such a component.  While many authors have elucidated methods of calculating the emission from small, thermally fluctuating grains (e.g. Draine \\& Andersen 1985; Dwek 1986; Leger, d'Hendecourt \\& D\\'{e}fourneau 1989; Guhathakurta \\& Draine 1989), their inclusion in radiative transfer calculations has been limited. Our extension of the \\dm model includes large grains, small grains ($a > 100$~\\AA\\ and $a \\le 100$~\\AA, respectively; see \\S\\ref{sec:trans}), and PAH molecules and we treat the heating and re-emission by each component in the appropriate regime.  In this paper, we present our model, concentrating on the dust heating and emission. Details of the Monte Carlo calculations are presented in a companion paper \\citep{gmwc00}. In \\S\\ref{sec:dust_model}, we discuss the details of our dust grain model; we review the relevant equations for determining the dust emission spectrum, describe our computational method, and discuss the details of the computation in \\S\\ref{sec:physics}; results of the model calculations are presented in \\S\\ref{sec:application} in the context of applications of \\dm to starburst galaxies, including a discussion of the response of the model SED to variations in the input parameters, e.g., the dust grain model, global and local geometries, heating sources (age, star formation rate [SFR]), size, and optical depth; we conclude with a summary and outline some future directions in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper and a companion paper \\citep{gmwc00}, we have presented the \\dm model, a self-consistent Monte Carlo radiative transfer and dust emission model.  The strengths of \\dm include:  \\begin{itemize} \\item Self-consistency.  No {\\em ad hoc} assumptions about the dust temperature are made; the dust is heated by the absorption of photons originating from sources included in the model.  The temperature distribution of the dust grains, and hence their emission spectrum, is calculated self-consistently by keeping track of the energy absorbed by dust in the Monte Carlo radiative transfer.  \\item Treatment of the heating of and emission from small grains, including the effects of temperature fluctuations resulting from the absorption of single photons.  \\item The ability to properly treat high optical depth situations. The iterative solution outlined above, wherein dust self-absorption is considered, permits the treatment of high optical depth cases, where the optically thin assumption may be violated even at IR wavelengths.  \\item Properly including the FIR emission from dust in \\dm provides additional information regarding the dust model, the nature of the heating source(s), and the physical size of the modeled region, not available through UV--NIR modeling alone.  \\item The use of Monte Carlo techniques to solve the radiative transfer equations allows the treatment of arbitrary dust and heating source distributions as well as inhomogeneous dust distributions (clumpiness).  Though we have concentrated on starburst galaxies in exploring the parameter space of \\dm$\\!\\!$, it is of general applicability and can be used to model e.g., dust tori in AGNs and Quasars, circumstellar discs, individual star forming regions, and reflection nebulae.  \\end{itemize}  To emphasize this last point, we point out that the fits to any UV--NIR SED are degenerate; there are in general several possible combinations of geometries, dust models and heating sources which can provide essentially identical fits to the SED.  However, by examining the IR dust spectrum (total energy emitted by dust, peak wavelength of the IR emission, the strength of the mid IR emission relative to longer IR wavelengths, and strength of IR absorption and emission features), some of the degeneracy in the UV--NIR models can be lifted.  Indeed, the FIR dust spectrum provides information on the physical size of the region not accessible to UV--NIR modeling alone.  Future papers will detail the application of \\dm to interpreting observations of astrophysical systems including individual starburst galaxies.  Improvements to the current model are also being implemented.  We have presented a fairly simple treatment of the PAH component here;  we have not included a means of varying the relative strengths of the PAH features.  Since the strengths of some of the features depend on the number of hydrogen atoms in the molecule, while others depend on the number of carbon atoms, the relative strengths can be varied by adjusting the hydrogen coverage, $x_H$ (see \\S\\ref{sec:dust_model}).  Applications to specific, individual objects will require this sort of fine tuning as the relative strengths of the MIR emission features are observed to change from object to object and environment to environment \\citep{cohen+89}.  A major area of effort will be in applying \\dm in a starburst model that will include multiple stellar populations and evolutionary effects.  This will require the inclusion of the effects of the interaction between on-going star formation and dust, including the evolution of the grains due to processing, formation, and destruction (e.g. Efstathiou, Rowan-Robinson \\& Siebenmorgen 2000)."
    },
    "0011/astro-ph0011095_arXiv.txt": {
        "abstract": "A primordial stochastic magnetic field will induce temperature and polarization fluctuations in the cosmic microwave background. We outline a calculation of the resulting fluctuation power spectra and present numerical results. ",
        "introduction": "The presence of a magnetic field in the early universe influences the evolution of metric perturbations, and as a result, might leave observable traces in the cosmic microwave background (CMB).  During the past few years the gravitational effects of a homogeneous magnetic field were presented by several authors; for a detailed review see \\cite{grasso} and references therein.  More realistic is the case of a stochastic magnetic field, because any causal generation mechanism results in stochastic fields.  Some estimates of CMB temperature power spectra from density perturbations induced by a stochastic magnetic field are given in \\cite{lemoine} and numerical simulations are presented in \\cite{koh}. The vorticity induced by a stochastic magnetic field was studied in \\cite{SB} and CMB temperature fluctuations were obtained in the case of a single $k$ mode. The gravitational waves generated by tangled magnetic fields and the resulting tensor CMB temperature fluctuations spectra are given in \\cite{DFK}.  Here we outline a comprehensive analytic calculation of temperature and polarization power spectra due to all types of metric perturbations arising from a primordial stochastic magnetic field, which we assume to be statistically homogeneous and isotropic and described by a power law. Detailed results will be presented elsewhere. ",
        "conclusions": ""
    },
    "0011/astro-ph0011430_arXiv.txt": {
        "abstract": "This contribution presents the optical spectral properties of FIR-luminous galaxies, whose distinctive feature is often the simultaneous presence in the spectra of a strong $\\rm H\\delta$ line in absorption and of emission lines (e(a) spectra). A discrepancy between the star formation rate estimated from the FIR luminosity and that derived from the $\\rm H\\alpha$ luminosity persists even after having corrected the $\\rm H\\alpha$ flux for dust according to the observed Balmer decrement. It is shown that the e(a) spectrum can be reproduced assuming a current starburst and a dust extinction affecting the youngest stellar populations much more than the older stars. ",
        "introduction": "The most distinctive feature in the optical spectra of FIR-luminous galaxies is not the equivalent width of the emission lines\\footnote{These are usually lower than those in HII galaxies and UV-bright starbursts.}, but the equivalent width of the $\\rm H\\delta$ line in absorption, which is generally stronger than in the spectra of optically selected galaxies. In fact, emission line spectra with EW($\\rm H\\delta)>4$ \\AA $\\,$ (e(a) spectra) are frequent among FIR galaxies, while the great majority of nearby spirals in optical samples have EW($\\rm H\\delta)<4$ \\AA $\\,$ (e(c) spectra) (Poggianti et al. 1999 [P99]; Poggianti \\& Wu 2000 [PW00]). The e(a) spectral class differs from the so-called \"k+a\" (or \"E+A\") galaxies which only display a strong $\\rm H\\delta$ line in absorption, but no emission lines.  What is the origin of the difference in the $\\rm H\\delta$ strength between dusty starburst galaxies and quiescent spirals? What star formation and dust properties generate the peculiar spectral combination found in e(a) galaxies? In the following I will address these issues presenting the analysis of a spectroscopic survey of luminous IRAS galaxies (PW00) and the results of an effort to model the e(a) spectrum in detail (Poggianti, Bressan \\& Franceschini 2001). ",
        "conclusions": ""
    },
    "0011/hep-ph0011333_arXiv.txt": {
        "abstract": "s{Important progress in understanding the behavior of hadronic matter at high density has been achieved recently, by adapting the techniques of condensed matter theory.  At asymptotic densities, the combination of asymptotic freedom and BCS theory make a rigorous analysis possible. New phases of matter with remarkable properties are predicted.  They provide a theoretical laboratory within which chiral symmetry breaking and confinement can be studied at weak coupling.  They may also play a role in the description of neutron star interiors. We discuss the phase diagram of QCD as a function of temperature and density, and close with a look at possible astrophysical signatures.}   \\vspace{1cm}  \\tableofcontents  \\newpage ",
        "introduction": "In this article we shall discuss the behavior of QCD at high density. We shall mainly be concerned, more precisely, with the regime of very large baryon number density and relatively low temperature.  The appeal of the subject should be obvious. It provides the answer to a child-like question: What happens to matter, as you squeeze it harder and harder?   Moreover, this regime may be realized in  neutron star interiors and in the violent events associated with collapse of massive stars or  collisions of neutron stars, so it is important for astrophysics.   Finally, we may hope to gain insight into QCD at moderate or low density --- the realm of ordinary matter and of terrestrial experiments --- by extrapolating from above.  What is less obvious, but turns out (with certain qualifications) to be true, is that the high density regime of QCD is tractable theoretically. Heuristic arguments to this effect, invoking asymptotic freedom,\\cite{AF} go back to the earliest days of the subject.\\cite{CollinsPerry} High density brings in a large energy scale, the chemical potential $\\mu$, and one might hope that the relevant coupling to describe the dynamics is $g(\\mu)$.  This becomes small as $\\mu \\rightarrow \\infty$, and suggests the possibility of a perturbative treatment.  This naive expectation, however, does not stand up to critical scrutiny.  As we shall discuss at length below, perturbation theory around the naive ground state  (free quark Fermi spheres) encounters infrared divergences.  Furthermore, the naive perturbative ground state is unstable.   Therefore, straightforward perturbative treatment of QCD at high density fails.  Fortunately, related difficulties have been met and overcome previously, in the theory of superconductivity.    There we learn that arbitrarily weak attractive interactions can change the ground state qualitatively.  In the true ground state there is an effective mass for photons --- the Meissner effect --- and energy gaps for charged excitations.   These phenomena remove potential infrared divergences, and render the perturbation theory around the true ground state regular (nondegenerate).   We can readily adapt  the methods of superconductivity theory to QCD.  It is also instructive to consider, in this connection, variants of QCD with different numbers of colors and flavors, and different spectra of quark masses.  A rich and highly structured theory emerges, displaying calculable, highly non-trivial dependence on all these variables.  The central result of the analysis is the identification of condensates in diquark channels, analogous to the Cooper pairs of electrons in ordinary superconductors. This is the phenomenon of color superconductivity.\\cite{Barrois,BailinLove,ARW1,RappETC}     Compared to ordinary superconductivity, color superconductivity, though it appears superficially to be more complex mathematically, is in a profound sense simpler and more directly related to fundamentals. Ordinary superconductivity takes place in solids and the accurate effective interactions are determined by band structure and other complicated effects. % Furthermore, ordinary superconductivity in a metal involves electron pairing, and the fundamental interaction between electrons (the screened Coulomb interaction) is repulsive. The effective attraction near the Fermi surface that leads to superconductivity arises in classic superconductors only as a subtle consequence of retarded phonon interactions, and in the cuprate superconductors through some mechanism yet unknown.  In color superconductivity, by contrast, the attractive interaction can arise already from the primary, strong, interactions. This has two consequences. First, the accurate form of these interactions can be calculated from first principles, using asymptotic freedom.  This makes calculations at high enough density robust. Second, at accessible densities, where the strong interactions are much stronger than the electromagnetic interactions, we expect the color superconductors themselves to be robust in the sense that the ratio of their gaps and critical temperatures to the Fermi energy should be quite large.   In QCD with three colors and three flavors, we find an improved ground state at high density, based on color superconductivity, around which weak-coupling perturbation theory is valid. In particular, all the colored degrees of freedom acquire gaps. Thus, the improved ground state differs qualitatively from the naive one.  The resulting predictions regarding the low-energy spectrum and dynamics are striking.  Color symmetry and chiral symmetry are spontaneously broken. The spectrum of elementary excitations is quite different from that found in naive perturbation theory.  Nominally massless quarks and gluons become massive, new massless collective modes appear, and various quantum numbers get modified.    All the elementary excitations carry integral electric charges.\\cite{CFL} Altogether, one finds an uncanny resemblance between the properties one computes at asymptotic densities, directly from the microscopic Lagrangian, and the properties one expects to hold at low density,  based on the known phenomenology of hadrons.    In particular, the traditional ``mysteries'' of confinement and chiral symmetry breaking are fully embodied in a controlled, fully microscopic, weak-coupling (but nonperturbative!) calculation, that accurately describes a directly physical, intrinsically interesting regime.\\cite{CFL,SW1}   Though some of the ideas go back a decade or more,\\cite{Barrois,BailinLove} the full power of color superconductivity to provide a rigorous foundation for the investigation of high density QCD has only become apparent relatively recently, and the subject is developing rapidly.  In this survey we shall emphasize what we see as the most fundamental ideas that have appeared in the field to date, and attempt to identify some significant challenges for the future.   Although we shall supply extensive references we will not attempt to catalogue all the very latest results, nor snapshot the developing state of the art in its technical and quantitative aspects.    In Section 2, we shall discuss in detail the high-density behavior of a slightly idealized version of real-world QCD, in which we imagine there are three flavors of massless quarks.  This is the case for which the clearest and most beautiful results emerge, based on the phenomenon of color-flavor locking. In Section 3, we shall discuss what changes must be made, in order to take into account the realistic quark spectrum.   We shall also briefly discuss QCD variants with unrealistic numbers of colors and flavors, which shed additional light on the theory. In Section 4, we shall discuss in greater depth the theoretical analysis which underlies the preceding sections, and the techniques it relies upon. In Section 5, we shall synthesize the preceding discussion, to produce a tentative sketch of the phase diagram for high-density QCD. Finally in Section 6 we shall discuss possible applications to astrophysics, including both existing results and some avenues that we believe might reward further investigation. ",
        "conclusions": ""
    },
    "0011/astro-ph0011160_arXiv.txt": {
        "abstract": "We investigate the star formation and dust extinction properties of very luminous infrared galaxies whose spectra display a strong $\\rm H\\delta$ line in absorption and a moderate [O{\\sc ii}] emission (e[a] spectrum). This spectral combination has been suggested to be a useful method to identify dusty starburst galaxies at any redshift on the basis of optical data alone. We compare the average e(a) optical spectrum with synthetic spectra that include both the stellar and the nebular contribution, allowing dust extinction to affect differentially the stellar populations of different ages. We find that reproducing the e(a) spectrum requires the youngest stellar generations to be significantly more extinguished by dust than older stellar populations, and implies a strong ongoing star formation activity at a level higher than in quiescent spirals. A model fitting the optical spectrum does not necessarily produce the observed FIR luminosity and this can be explained by the existence of stellar populations which are practically obscured at optical wavelengths. Models in which dust and stars are uniformly mixed yield a reddening of the emerging emission lines which is too low compared to observations: additional foreground reddening is required. ",
        "introduction": "The effects of dust extinction are crucial for interpreting the spectra of nearby and distant galaxies, particularly after that new facilities from the ground (e.g. {\\it SCUBA} on {\\it JCMT}) and from space ({\\it ISO}) have established that the IR emissivity of galaxies was much enhanced in the past with regard to what is typically observed in local galaxy populations (Elbaz et al. 1999; Barger et al. 2000, Franceschini et al. 2000, Ivison et al. 2000, Smail et al. 2000). Further evidence about the very important role played by dust reprocessing during the evolution of galaxies was provided by the detection of a diffuse far-IR background likely originating from distant and primeval galaxies, whose dominance over the optical background testifies hidden star formation at unexpectedly high rates (Puget et al. 1996; Dwek et al. 1998).  Waiting for future spectroscopic capabilities in the far-IR which will be offered by a variety of missions (SIRTF, FIRST, and NGST), only indirect clues about the relation between dust and stellar population emissions can be gained from optical spectra.  In the local universe, a large fraction of the FIR-luminous galaxies are characterized by a peculiar combination of spectral features in the optical: a strong $\\rm H\\delta$ line in absorption (EW$> 4$ \\AA) and a moderate [O{\\sc ii}] emission (EW$< 40$ \\AA) (Poggianti \\& Wu 2000, hereafter PW00). Galaxies with this type of spectra were named ``e(a)'' galaxies \\footnote{We stress that these are {\\it not} the so-called ``E+A'' (or k+a) spectra which \\it by definition \\rm have \\it no detectable emission lines \\rm (Dressler \\& Gunn 1983), at least at the level of the [O{\\sc ii}] detection limit of the spectroscopic surveys of distant galaxies (typically 3 \\AA). Nevertheless, it is possible that some of the k+a spectra are e(a) galaxies with an extremely low [O{\\sc ii}] (P99, Smail et al. 1999).} and were found to be quite numerous in the cluster and field environments at $z=0.4-0.5$ (Dressler et al. 1999, Poggianti et al. 1999, hereafter P99). The equivalent width of their $\\rm H\\delta$ line exceeds that of typical, quiescent spirals at low-z and their low [O{\\sc ii}]/$\\rm H\\alpha$ ratios are consistent with the emission line fluxes being highly extincted by dust (P99, PW00).  Interestingly, the combination of moderate emission lines and strong early Balmer lines in absorption has been subsequently found also in the first spectroscopic surveys of the high-redshift galaxy populations detected by {\\it ISO} (Aussel 1999, Flores et al. 1999). The association between the e(a) signature and the FIR-luminous galaxies seems to be confirmed at any redshift and suggests that the e(a) spectra -- despite their moderate emission lines -- actually belong to highly obscured, {\\it starburst} galaxies. One of the most important pieces of information that is still missing is the physical origin of this type of spectra: what star formation and dust properties produce this peculiar spectral combination? P99 and PW00 proposed that the e(a) spectrum originates from ``selective extinction'' in a galaxy where the youngest stars in the HII regions are much more dust-embedded than the older stellar populations, which are more widely distributed throughout the galaxy. Such an age-dependent obscuration would cause the emission lines to be highly extincted (hence producing the low blue-to-red emission line ratios observed) and would allow the older, A-type stars to dominate the integrated spectrum at $\\lambda \\sim 4000$ \\AA $\\,$ (thus explaining the strong $\\rm H\\delta$ line in absorption).  The existence of a selective extinction in galaxies is well known, but it is usually ignored when dealing with the global effects of dust on the spectral energy distribution of galaxies. Stars are expected to spend the very beginning of their evolution deeply embedded in dusty environments, later drifting away from or dispersing the molecular clouds where they were born (Calzetti et al. 1994 and references therein, Granato et al. 2000). As discussed in PW00, an age-dependent dust obscuration could explain the numerous observational studies which measure discrepant extinction values when using different spectral ranges/features (Israel \\& Kennicutt 1980, Viallefond et al. 1982, Caplan \\& Deharveng 1986, Fanelli et al. 1988, Bohlin et al. 1990, Keel 1993, Calzetti et al. 1994, Veilleux et al. 1995, Lancon et al. 1996, Mayya \\& Prabhu 1996, Calzetti 1997,1998, Mas-Hesse \\& Kunth 1999). Moreover, Zaritsky (1999) found a clear trend of decreasing mean extinction with declining stellar effective temperature (i.e. stellar age) in the Large Magellanic Cloud.  So far, no attempt has been made to model the whole e(a) spectrum in a quantitative way. The aim of this work is to test whether the average e(a) spectrum can be modelled as a current starburst assuming a selective dust extinction that increases as the age of the stellar population decreases. In the following we employ a spectrophotometric model that includes both the stellar component and the thermal emission of the gas ionized by young stars. Such a model will be considered successful only if it is able to reproduce simultaneously the shape of the optical observed continuum (3600-6900 \\AA) and the strength of the most important emission/absorption lines: [O{\\sc ii}]$\\lambda$3727, $\\rm H\\delta$ at $\\lambda = 4101$ \\AA, $\\rm H\\beta$ at $4862$ \\AA $\\,$ and $\\rm H\\alpha$ at $6563$ \\AA.  The plan of the paper is as follows: in \\S2 we present the spectroscopic catalog and the average observed e(a) spectrum and in \\S3 we describe the ingredients and the structure of our model. The results are discussed in \\S4, where we examine starburst models with selective extinction (\\S4.1), post-starburst models (\\S4.2) and models with constant extinction (\\S4.3). Our main conclusions are summarized in \\S5. ",
        "conclusions": "Spectra with a strong $\\rm H\\delta$ line in absorption and a moderate [O{\\sc ii}] line in emission -- named e(a) spectra -- are numerous among luminous infrared galaxies. We have shown that the optical e(a) spectrum can be well reproduced by a spectrophotometric model that includes both the stellar component and the emission lines and continuum of the ionized gas, {\\it assuming dust extinction to vary with the age of the stellar population}.  The main conclusions of this work can be summarized as follows:  1) On the basis of our models, two conditions are necessary to explain the e(a) spectra:  a) the youngest stellar generations (age $\\leq 10^7$) in e(a) galaxies must be highly obscured by dust. Their extinction values are significantly higher than those applying to the previous stellar populations. Models with a constant value of extinction are unsuccessful in fitting the optical spectrum;  b) these galaxies must host strong ongoing star formation at a level higher than in quiescent spirals. Post-starburst models with a small residual star formation fail to explain simultaneously the strength of the $\\rm H\\alpha$ and of the $\\rm H\\beta$ lines.  2) The degeneracy of the optical spectrum hinders a quantitative estimate of the burst duration and intensity. The spectrum is consistent with, but does not require, a ``long'' starburst that began $10^8-10^9$ yr ago.  3) A model fitting the e(a) spectrum does not necessarily produce the FIR/optical ratio observed: the optical spectrum (including the observed $\\rm H\\alpha$ luminosity and $\\rm H\\alpha/H\\beta$ ratio) does not constrain the current star formation rate. The observed FIR/V can be explained by starburst models if a significant fraction of the FIR luminosity originates in regions that are practically obscured at optical wavelengths.  4) All the results listed above refer to the case of a dust screen placed in front of each stellar generation. Models in which dust and stars are uniformly mixed fail to fit the e(a) spectrum, producing too low a reddening of the emerging emission lines. The observed reddening seems to be at least partly due to foreground dust.  5) The extinction correction factor at $\\rm H\\alpha$ found from the Balmer decrement in the e(a) Very Luminous Infrared galaxies of our sample ($E(\\rm H\\alpha)\\sim 15$, corresponding to a median E(B-V)=1.1, W98a) is in agreement with that derived with the same method in other FIR luminous samples (Kim et a. 1998, Veilleux et al. 1995, median E(B-V) for non-AGN galaxies =1.13 and 1.05, respectively). This factor is about 6 times higher than the average $E(\\rm H\\alpha)=2.5$ in nearby spirals (Kennicutt 1992). Ignoring the fact that the slit might have missed a fraction of the $\\rm H\\alpha$ luminosity and applying the extinction correction to derive the total current star formation rate from the relation $\\rm SFR(M_{\\odot} \\, yr^{-1}) = 0.8 \\times 10^{-41} \\, E(\\rm H\\alpha)) \\times L(H\\alpha) \\, erg \\, s^{-1} $ (Kennicutt 1992)\\footnote{For a Salpeter IMF (x=2.35, $0.1-100 M_{\\odot}$) and a standard Galactic extinction law ($R_V=3.1$).}, yields a value of the SFR that is still underestimated by a factor $\\sim 3$ compared to the SFR derived from the FIR luminosity."
    },
    "0011/astro-ph0011483_arXiv.txt": {
        "abstract": "Data taken with the HEGRA system of imaging atmospheric Cherenkov telescopes during the 1997 flares of Markarian 501 (Mkn 501) are reanalyzed using an algorithm providing improved energy resolution. A resolution of 10\\% to 12\\% is obtained by accounting for the variation of the Cherenkov light yield with the height of the shower maximum in the atmosphere. The improved energy resolution is particularly relevant for the study of the high-energy cutoff in the spectrum, which might be caused by interactions with the intergalactic infrared background radiation. The reanalysis presented here confirms the results obtained in the previous analysis, but hints a steeper slope of the spectrum in the region around 20~TeV. ",
        "introduction": "The attenuation of extragalactic gamma rays in the TeV energy range due to interactions with the cosmic infrared background radiation (Nikishov 1962; Gould \\& Schr\\'eder 1967; Stecker et al. 1992) is under intense discussion (see, e.g., Funk et al. 1998; Stecker \\& de Jaeger 1998; Aharonian et al. 1999a; Coppi and Aharonian 1999; Konopelko et al. 1999a; Stecker 1999; Krawczynski et al. 2000; Vassiliev 2000; and references given therein), in particular since the energy spectra determined for the AGN Mkn 501 (Aharonian et al. 1999a,b,c; Samuelson et al. 1998; Djannati-Atai et al. 1999), obtained during its 1997 flares, show indications of absorption features, particularly evident in the high-energy HEGRA data of Aharonian et al. (1999a), in the following referred to as Paper I. \\begin{figure*}[htb] \\begin{center} \\mbox{ \\epsfxsize7.5cm \\epsffile{h2398f1a.eps}} \\mbox{ \\epsfxsize7.2cm \\epsffile{h2398f1b.eps}} \\label{fig_abs} \\caption{(a) Attenuation of TeV gamma rays from Mkn 501 due to interaction with the infrared background radiation, for different assumptions concerning the density of the infrared background. Dotted line: following recent data from Finkbeiner et al. (2000), as shown by the solid line of Fig. 1 in Protheroe \\& Meyer (2000). Dashed-dotted line: using the minimal flux consistent with the $1 \\sigma$ statistical and systematic error bars of Finkbeiner et al. (2000). Full line: earlier LCFDM model of Primack et al. (1999). All curves use are calculated for the redshift $z \\approx 0.03$ of Mkn 501, assuming a Hubble constant of 65 km s$^{-1}$Mpc$^{-1}$. (b) Spectral energy distribution $E^2$d$N$/d$E$ of TeV gamma rays from Mkn 501 at the source, reconstructed by dividing the measured distribution (HEGRA, Paper I, shown as triangles) by the attenuation. Filled circles: using the attenuation indicated by the dotted curve in part (a) of this Figure. Open circles: using the attenuation indicated by the dashed-dotted curve, shifting the HEGRA data down in energy by 15\\% -- the maximal systematic error on the energy scale -- and including the systematic error bars of the HEGRA measurement.} \\end{center} \\end{figure*} The level of the cosmic infrared background is related to the evolution of galaxies and can serve to distinguish different cosmological models (see, e.g., Primack et al. 1999). However, due to the intense foreground radiation, a reliable direct measurement of the infrared background is a difficult task. Gamma rays in the TeV energy range provide an independent and complementary approach to study the infrared background radiation. Since the $\\gamma \\gamma \\rightarrow e^+e^-$ interaction cross section peaks near threshold, the absorption of gamma rays of energy $E$ probes the background level around a wavelength $\\lambda/\\mu\\mbox{m} \\approx \\kappa E/\\mbox{TeV}$, where $\\kappa$ varies between 1 and 2 depending on the spectral shape of the background radiation. There is almost a one-to-one correspondence between the energy-dependent attenuation and the spectral density of the infrared background radiation. Recent analyses by Finkbeiner et al. (2000) of the background level at long wavelengths, at 60~$\\mu$m and 100~$\\mu$m, show a significantly higher flux than expected, which results in strong absorption of Mkn 501 gamma rays above 10 TeV (Fig.~\\ref{fig_abs}(a)). While earlier analyses using a lower level of the infrared background (e.g., the models of Primack et al. (1999), also illustrated in Fig.~\\ref{fig_abs}(a)) found good consistency with the TeV data (see, e.g., Konopelko et al. (1999a)), these new results appear to be in contradiction with measurements using the HEGRA Cherenkov telescopes, which show a gamma-ray flux from Mkn 501 up to at least 16~TeV (Paper I). In order to explain the observed flux in the presence of strong high-energy attenuation, the spectral distribution of gamma-rays at the source would have to exhibit a rise at high energies (Fig.~\\ref{fig_abs}(b)), inconsistent both with acceleration models and with energetics of the source (see, e.g., Paper I and Coppi and Aharonian 1999; Guy et al. 2000; Protheroe \\& Meyer 2000; Finkbeiner et al. 2000; and references given therein). Exotic solutions of this ``gamma-ray crisis'' have been proposed, such as violations of Lorentz invariance (Kifune 1999) or photon condensates (Harwit et al. 1999); the latter explanation has been questioned on rather general grounds (Levinson 2000) and seems already to be ruled out by a recent HEGRA analysis (Aharonian et al. 2000).  The TeV gamma-ray data and the infrared background data can be made marginally consistent by fully exploiting the systematic errors of both measurements, i.e., taking the minimal infrared flux values and gamma ray flux values and reducing the HEGRA energy scale by 15\\%. In this case, a marginally acceptable source spectrum is achieved, see Fig.~\\ref{fig_abs}(b), in particular given that the highest HEGRA data point is, at the $2\\sigma$ level, consistent with spill-over from the lower energy bins. The bulk of the change in the calculated source spectrum -- a factor 100 for the highest data point -- can be attributed to the uncertainty in the infrared data; the effect of the HEGRA uncertainties is smaller -- about a factor 10 for the highest point.  The interpretation of the experimental results is currently limited by the systematic errors both of the measurement of the infrared background, and of the TeV spectra. One potentially very important factor is the energy resolution in the detection of the TeV gamma rays. The HEGRA system of Cherenkov telescopes, which provides the highest-energy data, has an energy resolution around 20\\% (see Aharonian et al. 1999b, Konopelko et al. 1999b, Paper I). While this resolution is ample for the reconstruction of the usual power-law energy spectra, it may smear out the extremely sharp cutoff features expected for absorption due to infrared background, in particular for high long-wavelength infrared background levels (Fig.~\\ref{fig_abs}(a)). The highest-energy gamma rays detected could simply represent spill-over from lower energies due to the finite energy resolution. Indeed, while Aharonian et al. (Paper I) measure photons with reconstructed energies up to 21.5 TeV and beyond, they state that a sharp cutoff at 16 TeV cannot be excluded by the data. It is obviously of considerable interest to improve the energy reconstruction of the system of telescopes, in order to more precisely map a possible cutoff. A scheme to improve the energy resolution of systems of Cherenkov telescopes by roughly a factor 2 has been discussed by Hofmann et al. (2000), making use of the abundant information provided by such systems for each individual gamma-ray air shower. In this note, we present a reanalysis of the HEGRA Mkn 501 spectral data with such improved analysis algorithms. ",
        "conclusions": ""
    },
    "0011/astro-ph0011356_arXiv.txt": {
        "abstract": "The radiation observed by blazars is believed to originate from the transformation of bulk kinetic energy of relativistic jets into random energy. A simple way to achieve this is to have an intermittent central power source, producing shells of plasma with different bulk Lorentz factors. These shells will collide at some distance from the center, producing shocks and then radiation. This scenario, called {\\it internal shock model}, is thought to be at the origin of the $\\gamma$--rays observed in gamma--ray bursts and can work even better in blazars. It accounts for the observed key characteristics of these objects, including the fact that radiation must be preferentially produced at a few hundreds of Schwarzschild radii from the center, but continues to be produced all along the jet. At the kpc scale and beyond, the slowly moving parts of a (straight) jet can be illuminated by the beamed radiation of the core, while the fast parts of the jet will see enhanced cosmic microwave radiation. In both cases the Inverse Compton process can be the dominant radiation process, leading to a copious production of high energy (X--rays and beyond) radiation in both radio loud quasars and radio--galaxies. ",
        "introduction": "We believe that the radiation we see from blazars comes from the transformation of bulk kinetic into random energy of particles, which then produce beamed emission. How to produce the large velocities of the plasma in the jet and which is the dissipation mechanism are still a matter of debate, but there is no doubt that nature succeds in producing collimated outflows with bulk Lorentz factors $\\Gamma\\sim$ 5--20 for blazars, and even higher for gamma--ray bursts. Only in recent years we began to estimate the power of jets, through the radiation they emit (e.g. Celotti \\& Fabian, 1993) and especially through the energy required to be transported to the lobes (Rawling \\& Saunders, 1991). It has been found that the observed jet radiation must be a small fraction of the total energetics, even of the last decade witnessed a factor 10 increase in the power observed to be emitted by blazars as a class, thanks to the high energy $\\gamma$--ray observations of EGRET, onboard the Compton Gamma Ray Observatory satellite. \\begin{figure} \\plotone{ghisellini_fig1.ps} \\caption{Three examples of SED of blazars to illustrate the blazar sequence. The top panel shows the most distant radio--loud AGN known, at $z=4.72$ (from Fabian et al. 2000). Its luminosity in hard X--rays exceeds, if isotropic, $10^{49}$ erg s$^{-1}$. Note that the peak of the inverse Compton emission is in the MeV band, and that the hard X--ray emission dominates the power output. In the mid panel we show the intermediate BL Lac object ON 231, in which the synchrotron emission dominates the steep soft X--ray flux, and a very flat inverse Compton component dominates above a few keV (from Tagliaferri et al., 2000). In the bottom panel we show the SED of the extreme BL Lac 1426+428, in which the synchrotron component peaks above 100 keV (from Costamante et al., 2000). This is the third example of BL Lac object with a synchrotron peak located above 100 keV, besides Mkn 501 and 1ES 2344+514. In these low luminosity class of sources, the emitting electrons can attain the highest energies, making these objects the best candidates to be detected in the TeV band. Note the broad band X--ray range of {\\it Beppo}SAX and how useful it is to characterize the SED in all three classes of blazars. } \\end{figure} The EGRET observations, and the detection of a few sources (Mkn 421, Mkn 501, 2344+514 and PKS 2155--304) in the TeV band by ground based Cherenkov telescopes, renewed the interest about blazars, allowing the discovery that their Spectral Energy Distribution (SED) is characterized by two broad peaks, whose location is a function of the observed bolometric luminosity (Fossati et al. 1998, Ghisellini et al. 1998). These peaks have been interpreted as due to synchrotron and inverse Compton radiation, respectively. Blazars form a well defined sequence, with low powerful objects having both peaks at a similar level of luminosity, and located at higher frequencies than in more powerful objects, in which the inverse Compton peak dominates the emission. In Fig. 1 we show three examples of SED of blazars with different power, to illustrate the overall behavior and what can be the contribution of X--ray observations in these three classes of objects. Recent observations of high redshift ($z>4$) blazars (Celotti, these proceedings), and of low power BL Lacs (Costamante et al. these proceedings) have extended the blazar sequence at both ends, confirming the original trend.  At the high luminosity end of the sequence we find interesting lower limits on the bulk kinetic power that the jet can carry, requiring  it to be larger than the power dissipated in radiation, derived dividing the apparent luminosity (assuming isotropy and no beaming) by the square of the bulk Lorentz factor. Results indicate that jet of FSRQ (flat spectrum radio quasars) must have a large kinetic power (Celotti, these proceedings). As an example, PKS 0836+710 has an apparent luminosity of $10^{49}$ erg s$^{-1}$, which requires a jet power of at least $10^{47}/\\Gamma_1^2$ erg s$^{-1}$ (Tavecchio et al., 2000a). Note that in these sources most of the jet power must not be dissipated through radiation, but must feed the extended radio structures.  At the low luminosity end of the blazar sequence we find objects whose synchrotron spectrum peaks in the X--ray band, indicating very large energies of the emitting electrons. Here we can learn about the acceleration mechanism, and find good candidates to be detected in the TeV band (Costamante et al, these proceedings).  Here I will focus on two main topics, namely how the {\\it internal shock scenario} can explain the main characteristics of blazars, and how the large (and very large) scale X--ray jets recently observed by $Chandra$ can be interpreted. ",
        "conclusions": "Internal shocks can dissipate the right amount of jet power at the right location, originating in a natural way the violent variability observed in blazars. This scenario has the virtue to be at the same time simple and quantitative, offering a coherent view of almost all the radiative jet, from milliparsecs to kiloparsecs. In the $\\gamma$--ray band we expect the most pronounced variability, in agreement with observations, and we are now working to find how the $\\gamma$--ray duty cycle (i.e. the fraction of the time spent in high $\\gamma$--ray states) scales with the initial separation of the shells and their initial bulk Lorentz factor to compare it with data coming from the foreseen $\\gamma$--ray satellites AGILE and GLAST.  We hope that this scenario will also explain the observed ``blazar sequence\", linking the overall blazar spectrum with the jet power (i.e. the bulk motion power). We also hope, with longer simulations, to be able to see if there is a relation between the strongest $\\gamma$--ray flares and the birth of radio superluminal blobs at the pc scale.  The emission predicted by the internal shock scenario (as it is now) beyond the kpc scale likely underestimates the X--ray flux (both soft and hard). At these distances the energy densities in magnetic field and locally produced synchrotron radiation are very small, and seed photons of the cosmic microwave background on one hand and of the core of the jet on the other hand are important contributors to the inverse scattering process for fast spines and slow layers, respectively."
    },
    "0011/astro-ph0011026_arXiv.txt": {
        "abstract": "{ We present the first high resolution X-ray spectrum of the bright O4Ief supergiant star $\\zeta$~Puppis, obtained with the Reflection Grating Spectrometer on-board \\xmm.  The spectrum exhibits bright emission lines of hydrogen-like and helium-like ions of nitrogen, oxygen, neon, magnesium, and silicon, as well as neon-like ions of iron.  The lines are all significantly resolved, with characteristic velocity widths of order $1000\\, -\\, 1500\\, {\\rm km\\, s^{-1}}$. The nitrogen lines are especially strong, and indicate that the shocked gas in the wind is mixed with CNO-burned material, as has been previously inferred for the atmosphere of this star from ultraviolet spectra. We find that the forbidden to intercombination line ratios within the helium-like triplets are anomalously low for N VI, O VII, and Ne IX.  While this is sometimes indicative of high electron density, we show that in this case, it is instead caused by the intense ultraviolet radiation field of the star.  We use this interpretation to derive constraints on the location of the X-ray emitting shocks within the wind that are consistent with current theoretical models for this system. ",
        "introduction": "The initial discovery of X-ray emission from O stars with the {\\it Einstein\\ Observatory} in the late 1970s (Harnden et al. \\cite{HarnEL}), sparked a vigorous field of research aimed at better understanding the production of hot gas in such systems.  Most current models invoke hydrodynamic shocks resulting from intrinsic instabilities in the massive, radiatively driven winds from these stars (for a detailed review see Feldmeier et al. \\cite{Feld}).  Support for this picture comes from the fact that the X-ray flux is not highly absorbed at low energies (Cassinelli \\& Swank \\cite{CS}; Corcoran et al. \\cite{CorEL}), which is expected if the X-ray emitting gas is distributed throughout the wind rather than in some form of hot corona in the outer atmosphere of the star.  However, the X-ray observations to date have not been especially constraining for stellar wind models.  This is primarily due to the low spectral resolution of the available nondispersive detectors, which has precluded the study of individual atomic features so crucial to the unambiguous determination of physical conditions in the shocked gas.  In this Letter, we present one of the first high resolution X-ray spectra of an early-type star, the O4Ief supergiant $\\zeta$~Puppis, which we have obtained with the Reflection Grating Spectrometer (RGS) experiment on the \\xmm\\ Observatory.  $\\zeta$~Pup is an excellent target for such work since it is the brightest O star in the sky, and has consequently been very well studied at longer wavelengths.  In particular, Pauldrach et al. (\\cite{PaulEL}) have developed a very detailed NLTE model of the atmosphere and wind of this star, constrained by high quality ultraviolet spectra from Copernicus and IUE.  They find a mass loss rate of $\\sim 5.1 \\times 10^{-6}\\, {\\rm M_{\\odot}\\, yr^{-1}}$ and an effective temperature $\\sim 42\\, 000\\, {\\rm K}$.  The derived abundances indicate that the atmosphere is mixed with CNO-processed material, consistent with the theoretical picture of a highly evolved star near the end of core hydrogen burning. Hillier et al. (\\cite{HilEL}) found that the Pauldrach et al. model is consistent with the ROSAT X-ray spectrum of $\\zeta$~Pup, but only if the X-ray emission arises in shocks distributed throughout the wind.  At the very lowest energies ($\\leq 200\\, {\\rm eV}$), they find that the emitting material must be $\\geq 100 \\rm R_*$ away from the star.  This is a consequence of the fact that in the Pauldrach et al. model, helium is mostly only singly ionized in the outer regions of the wind, so that the photoelectric opacity is very high at soft X-ray energies.  Our \\xmm\\ RGS spectrum is dominated by broad emission lines of mostly hydrogen-like and helium-like charge states of nitrogen, oxygen, neon, magnesium, and silicon, and neon-like ions of iron.  The data provide a number of important constraints on the nature and location of the X-ray emitting material in the $\\zeta$~Pup wind.  Of particular interest is the fact that we see a suppression of the forbidden line and enhancement of the intercombination line in the helium-like triplets of nitrogen, oxygen, and neon.  We show that this is a natural consequence of the intense ultraviolet radiation field in the wind, and that it allows us to place constraints on the location of the X-ray emitting shocks relative to the star.  Additional constraints come from the emission line velocity profiles.  We show that the data are consistent with the predictions of the Pauldrach et al. model.  In section 2, we describe the details of the RGS observation, the nature of our data reduction and analysis, and the key observational features of the spectrum.  In section 3, we consider the implications of our results for our understanding of $\\zeta$~Pup, and of the X-ray emission from O stars in general. ",
        "conclusions": "The RGS spectrum we have presented provides very strong confirmation for a number of key aspects of conventional models of X-ray emission from early-type stars in general, and for $\\zeta$~Pup in particular.  The emission is dominated by broad emission lines indicative of hot plasma that is flowing outward with the wind.  We do not see evidence for strong attenuation at low energies, which confirms that the X-ray emitting regions are not concentrated at the base of the wind, but instead are distributed out to large radii. These very simple results agree well with the predictions of wind instability models.  As we have shown, the suppression of the forbidden lines in the He-like triplets of low Z elements allows us to derive upper limits to the radii of the emitting shocks in each case.  It is interesting to compare these values with Fig. 1 of Hillier et al. (\\cite{HilEL}), which is a plot of the radius at which optical depth unity is reached in the wind as a function of X-ray energy for the NLTE model of $\\zeta$~Pup calculated by Pauldrach et al. (\\cite{PaulEL}).  At the energies of the N VI, O VII, Ne IX, Mg XI, and Si XIII lines, unit optical depth is achieved at 22, 23, 9, 3.5, and 2.5 stellar radii, respectively, in this model.  These values are quite compatible with both our derived upper limits for N VI, O VII, Ne IX, and Mg XI, and our derived lower limit for Si XIII. Since the density in the wind drops off like $r^{-2}$, and the emissivity is proportional to $n^{2}$, we expect radii characteristic of the smallest radius at which the overlying wind is still transparent to the respective line, in each case. Our results are consistent with this expectation.  Further support for this picture comes from the observed velocity profiles. The higher Z lines have characteristic widths $\\sim \\, 1000\\, {\\rm km\\, s^{-1}}$, whereas the N VII line is distinctly broader.  The terminal velocity in the $\\zeta$~Pup wind is $2260\\, {\\rm km\\, s^{-1}}$ (Groenewegen et al. \\cite{GroeEL}), so the higher Z lines are most likely emitted at only a few stellar radii, whereas N VII can come from considerably further out.  We have shown that the respective line intensities are suggestive of significant enhancements of the nitrogen abundance relative to carbon and oxygen, as one would expect for CNO-processed material.  This, again, agrees well with the Pauldrach et al. model. Meynet \\& Maeder (\\cite{MeMa}) recently presented new evolutionary models for rotating single stars. They found that rotational mixing produces a significant surface helium and nitrogen enhancement. Meynet \\& Maeder suggested that stars with enhanced He-abundances and large projected rotational velocities are natural descendants of very fast rotating main-sequence stars. In these stars, the chemical enrichment at the surface is very fast and as a result of the strong rotational mixing the chemical structure of these stars could be near homogeneity. With its large projected rotational velocity of $v\\,\\sin{i} \\simeq 203$\\,km\\,s$^{-1}$ (Penny \\cite{Pen}), $\\zeta$\\,Pup most probably falls into this category."
    },
    "0011/astro-ph0011210_arXiv.txt": {
        "abstract": "We examine the Butcher-Oemler effect and its cluster richness dependence in the largest sample studied to date: 295 Abell clusters. We find a strong correlation between cluster richness and the fraction of blue galaxies, \\fb, at every redshift.  The slope of the $f_B(z)$ relation is similar for all richnesses, but at a given redshift, \\fb\\ is systematically higher for poor clusters.  This is the chief cause of scatter in the \\fb\\ vs. $z$ diagram: the spread caused by the richness dependence is comparable to the trend in \\fb\\ over a typical redshift baseline, so that conclusions drawn from smaller samples have varied widely.  The two parameters, $z$, and a consistently defined projected galaxy number density, $N$, together account for all of the observed variation in \\fb\\ within the measurement errors. The redshift evolution of \\fb\\ is real, and occurs at approximately the same rate for clusters of all richness classes. ",
        "introduction": "The Butcher-Oemler (BO) effect provided some of the first evidence for the evolution of galaxies and clusters.  \\BOa (1978, 1984) found an excess of blue galaxies in high redshift clusters in comparison with the typical early-type population observed in the central region of local clusters (Dressler 1980). The \\BOe has also been observed in more recent studies, which indicate an even stronger evolution of the fraction of blue galaxies in clusters (Rakos \\& Schombert 1995, Margoniner \\& de Carvalho 2000).  The BO effect is an indicator of evolution in the cluster galaxy population, which may result from changes in the morphology and star-formation rates of member galaxies with redshift.  The fact that blue galaxies are more commonly found at higher redshifts, and the observation of an apparent excess of S0 galaxies in low redshift clusters, lead to the suggestion by Larson \\etal (1980) of an evolutionary connection between S0 and spiral galaxies.  This idea can explain the population of blue galaxies observed by Butcher \\& Oemler as spiral galaxies seen just before running out of gas, and the disappearance of this population in more evolved, low redshift rich clusters (Dressler \\etal 1997, Couch \\etal 1998).  In the last decade, high-resolution \\HST images have allowed the determination of the morphology of these high redshift blue galaxies.  Dressler \\etal (1994), Couch \\etal (1994, 1998), and Oemler \\etal (1997) found that the \\BOa galaxies are predominantly normal late-type spirals, and that dynamical interactions and mergers between galaxies may be a secondary process responsible for the enhanced star formation.  A recent study by Fasano \\etal (2000) indicates that as the redshift decreases, the S0 cluster population increases while the number of spiral galaxies decreases, supporting Larson's original idea of spirals evolving into S0s.  A second factor affecting the number of blue galaxies may be their environment. The fraction of blue galaxies seems to be correlated with local galaxy density, and with the degree of substructure in the cluster.  Many studies have shown that the blue galaxies lie preferentially in the outer, lower density cluster regions (Butcher \\& Oemler 1984, Rakos \\etal 1997, Margoniner \\& de Carvalho 2000).  Also, Smail \\etal (1998) studied 10 massive clusters and found that the fraction of blue galaxies in these clusters is smaller than observed in \\BOa (1984) clusters at the same redshift range.  However, an opposite trend is found in non-relaxed clusters with a high degree of substructure, where the fraction of blue galaxies is higher than that observed in regular clusters (Caldwell \\& Rose 1997, Metevier \\etal 2000).  The evolution of these galaxies is probably correlated with its environment in the sense that spirals might consume and/or lose their gas, evolving to SOs, while falling into the higher density cluster regions.  Although the exact mechanism responsible for the \\BOe is not completely understood (Kodama \\& Bower 2000), most authors seem to agree that the effect is real.  A different idea is presented by Andreon \\& Ettori (1999) who show that the mean X-ray luminosity and richness of \\BOa (1984) clusters increases with redshift, and argue that the increase in the fraction of blue galaxies with redshift may not represent the evolution of a single class of objects.  We present observations of the \\BOe in a large sample of 295 Abell (1958, 1989) clusters of all richnesses, with no further selection on the basis of morphology or X-ray luminosities. This is important because all previous studies were based on samples biased toward richer clusters, with some samples being further selected according to morphology and/or X-ray luminosities, so that a combination of different selection effects could be mimicking the observed $f_B(z)$ relation.  Although our sample inherits biases existent in the Abell catalog, it should contain a more representative variety of clusters (in terms of degree of substructure, richness, and mass) at each redshift than any previous sample, and because of its large size should allow a better determination of the relation.  Any results \u0018driven by selection effects present in previous samples might become apparent when compared with this one.  We describe the data in Section 2, the \\BOe analysis in Section 3, and our conclusions in Section 4. ",
        "conclusions": "We compute \\fb\\ for 295 randomly selected Abell (1989) clusters, including galaxies as faint as $M^* + 2$ to sample a larger range of the luminosity function and provide better statistics in each cluster. The resulting \\fb\\ shows a stronger trend with redshift than did \\fb\\ as originally defined by BO84, consistent with the idea that the \\BOe is stronger among the late-type spirals and irregulars which dominate the galaxy populations at intermediate and lower luminosities.  This is the first sample large enough to allow the study of the \\fb\\ variations at a given redshift. A $\\chi_{\\nu}$ test applied to a simple linear $f_B(z)$ relation shows that the scatter in \\fb\\ at a given redshift is consistently larger than the measurement error, indicating a real cluster-to-cluster variation.  We investigate the richness dependence of the Butcher-Oemler effect, and find a strong correlation between \\fb\\ and galaxy counts in the sense that poorer clusters tend to have larger \\fb\\ than richer clusters at the same redshift.  The inclusion of poor clusters tends therefore to increase the slope of the $f_B(z)$ relation, and is another factor (together with the inclusion of fainter galaxies) responsible for the stronger evolution observed in this sample when compared to previous works based mostly on rich clusters.  We show that the fraction of blue galaxies in a cluster can be completely determined by its redshift and richness (within the measurement errors). The evolution in \\fb\\ with redshift is real, and occurs at approximately the same rate for clusters of all richnesses."
    },
    "0011/astro-ph0011282_arXiv.txt": {
        "abstract": "It has been suggested that circumstellar disks evolve from dense, actively accreting structures to low-mass, replenished remnants. During this transition, grains may assemble into planetesimals, or the disk may be cleared by newborn planets. Recently identified nearby groups of young stars provide valuable laboratories for probing disk evolution. I discuss the properties of dust disks in the TW Hydrae Association and the MBM 12 cloud, and compare the results to other studies of disk evolution and planet formation timescales. ",
        "introduction": "Planetary systems form out of circumstellar disks that are the remnants of star formation (Shu, Adams, \\& Lizano 1987). The young T Tauri stars with ages of $\\sim$1 Myr, which frequently have optically-thick, actively-accreting disks (Strom, Edwards \\& Skrutskie 1993), are thought to represent a stage prior to the main epoch of planet formation. In contrast, the dust disks of the so-called Vega-like objects, main sequence stars whose ages could be as large as 1 Gyr, appear to be more evolved; most of the dust has presumably coagulated into planets or planetesimals, and the remaining dust in the optically-thin disk is continually replenished by collisions between larger bodies (Lagrange, Backman \\& Artymowicz 2000). However, the age at which this transition occurs is not well constrained. One reason is a paucity of intermediate-age ($t \\sim$10 Myr) nearby stars in current samples. Unfortunately, the age estimates of early-type isolated main sequence stars are highly uncertain.  The two competing theories for making gas giant planets in disks predict very different timescales. Core accretion mechanism, in which a solid core builds up from icy and rocky planetesimals and later acquires an envelope of nebular gas, is expected to take $\\sim$10 Myr (e.g., Pollack et al. 1996). On the other hand, the disk instability mechanism, where a gravitationally unstable disk fragments directly into self-gravitating protoplanetary clumps, is a much faster process, forming giant planets in $\\sim$0.1 Myr (e.g., Boss 2000). Therefore, better observational constraints on disk evolution timescales may allow us to distinguish between the two scenarios. ",
        "conclusions": ""
    },
    "0011/astro-ph0011557_arXiv.txt": {
        "abstract": "\\index{abstract} We outline here the next generation of cluster--finding algorithms. We show how advances in Computer Science and Statistics have helped develop robust, fast algorithms for finding clusters of galaxies in large multi--dimensional astronomical databases like the Sloan Digital Sky Survey (SDSS). Specifically, this paper presents four new advances: (1) A new semi-parametric algorithm -- nicknamed ``C4'' -- for jointly finding clusters of galaxies in the SDSS and ROSAT All--Sky Survey databases; (2) The introduction of the {\\it False Discovery Rate} into Astronomy; (3) The role of kernel shape in optimizing cluster detection; (4) A new determination of the X--ray Cluster Luminosity Function which has bearing on the existence of a ``deficit'' of high redshift, high luminosity clusters. This research is part of our ``Computational AstroStatistics'' collaboration (see Nichol et al. 2000) and the algorithms and techniques discussed herein will form part of the ``Virtual Observatory'' analysis toolkit. ",
        "introduction": "Clusters of galaxies are critical cosmological probes for two fundamental reasons. First, galaxy clusters are the most massive virialised objects in the universe and reside in the tail of the mass distribution function. Thus the distribution, density and properties of clusters are very sensitive to the mean mass density of the universe. For example, in a low $\\Omega_m$ universe, the evolution of the mass function terminates at early epochs, thus ``freezing'' the number of massive systems. The opposite is true for $\\Omega_m=1$, were the mass function continues to evolve up to the present epoch (see Press \\& Schtecter 1974; PS).  Unlike other cluster methods of measuring $\\Omega_m$ (mass-to--light ratios, baryon fractions {\\it etc}), this technique provides a global measure of $\\Omega_m$ and is relatively insensitive to $\\Omega_{\\Lambda}$ (see Romer et al. 2000). For a more complete description of PS and this effect, the reader is referred to Viana \\& Liddle (1999), Reichart et al. (1999) \\& Borgani et al. (1999), and references therein.  Many authors have used this technique with a preponderance of the evidence for a low $\\Omega_m$, but the observed scatter between the various measurements is large.  For example, we have seen $\\Omega_m=0.3\\pm0.1$ (Bahcall et al. 1997), $\\Omega_m=0.45\\pm0.2$ (Eke et al. 1998), $\\Omega_m=0.5\\pm0.14$ (Henry 1997), $\\Omega_m\\sim0.75$ (Viana \\& Liddle 1999), $\\Omega_m\\sim0.85\\pm0.2$ (Sadat et al. 1998), and $\\Omega=0.96^{+0.36}_{-0.32}$ (Reichart et al. 1999). This scatter could be the combination of several effects including small sample sizes, errors in the survey selection functions as well as the necessity to compare local samples of clusters to more distant samples in any effort to see an evolutionary signal.   Secondly, each cluster represents a sample of galaxies that has formed at roughly the same time under roughly the same initial conditions. Thus, each cluster is a laboratory for understanding how galaxies form, evolve and interact with their environment. Such phenomena are usually known as the {\\em morphology-density relation} and/or the {\\em Butcher--Oemler Effect}.  There are many proposed models to explain these phenomena including ram--pressure stripping, pressure--induced star formation, temperature--inhibited star formation, galaxy harassment, shocks from cluster merger events {\\it etc.} As yet, we do not have a definitive answer but observations of clusters - as a function of the cluster properties - will help to dis--entangle these different possible mechanisms.  It is clear that we need larger, more objective, cluster catalogs to help obtain high precision measurements of $\\Omega_m$ as well as determine the physical mechanism(s) behind galaxy evolution in clusters (see Kron 1993 for a discussion of the role selection effects could play in the Butcher--Oemler Effect).  To significantly move beyond previous work, any new sample of clusters should possess the following key attributes: {\\it i)} Contain many thousands of clusters; {\\it ii)} Possess a well understood selection function; {\\it iii)} Probe a broad range in redshift and mass, and {\\it iv)} Possess physically meaningful cluster parameters.  There are several cluster projects under-way that satisfy these criteria (see, for example, Collins et al. 2000; Romer et al. 2000; Gladder \\& Yee 2000; Ebeling et al. 2000) including the SDSS--RASS Cluster (SRC) survey which we discuss below. ",
        "conclusions": ""
    },
    "0011/astro-ph0011411_arXiv.txt": {
        "abstract": "High quality observations of the low degree $p$-modes exist for almost two complete solar cycles using the solar spectrophotometer Mark-I, located and operating at the Observatorio del Teide (Tenerife, Spain). In this work, the observations available have been re-analyzed over a much wider time interval than before. We analyze the time variation of the yearly frequency shift and its frequency dependence. This information will be used in order to average annual power spectra by removing the effect of the solar cycle. Using this average power spectrum, a new estimate of the rotational splittings is attempted. ",
        "introduction": "Understanding the well known solar variability is one of the major goals of  solar physics. In the last years, the helioseismic data, by means of the frequency shifts of solar $p$-modes, have demonstrated to be very sensitive to the solar activity cycle. In particular, the existence of a positive shift of $\\sim$0.4 $\\mu$Hz, peak-to-peak,  with the solar cycle is very well known (R\\'egulo {\\it et al.} 1994, Jim\\'enez-Reyes {\\it et al.} 1998).  The spectrophotomer Mark-I, has been collecting solar observations for almost two complete solar cycles. The available database, one of the longest for low degrees, could be used to measure the mode parameters. However, the variation of the central frequency is an important effect which should be taken into account in order to use long time series.  In this work, we have analyzed the time variation of the frequency shift and its frequency dependence. Using this information we are able to average the annual spectra which is then fitted in order to measure the mode parameters, specially the rotational splittings. ",
        "conclusions": "An analysis of the $p$-mode frequency shift over more than one solar cycle has been carried out. The time dependence of the frequency shift as well as its frequency dependence have been parameterized. Then, this information has been used in order to shift each one of the annual spectra, removing therefore the signature of the solar activity cycle.  The rotational splittings are in good agreement with those obtained by Bertello {\\it et al.} (2000). Thus, the rotation rate is very close to 435 nHz like the rest of the radiative part. However, there is still some fluctuation at lower turning point. The comparison with the splittings measured by Chaplin {\\it et al.} (1999) shows a significant difference.  The method presented here can be useful in the analysis of long time series, covering an important part of the solar cycle. Last but not least, if all the ground instruments (e.g. BiSON, IRIS, ECHO) are joined, then the quality and also the duty cycle of the time series will increase, reducing the systematic errors, allowing new perspectives in the analysis of the solar activity cycle using the frequency shift as a new solar index."
    },
    "0011/astro-ph0011377_arXiv.txt": {
        "abstract": "The launch of the Chandra X-ray Observatory in July 1999 opened a new era in X-ray astronomy. Its unprecedented, $<0.5$\" spatial resolution and low background are providing views of the X-ray sky 10-100 times fainter than previously possible. We have initiated a serendipitous survey (ChaMP) using Chandra archival data to flux limits covering the range between those reached by current satellites and those of the small area Chandra deep surveys. We estimate the survey will cover $\\sim 5$ sq.deg. per year to X-ray fluxes (2$-$10 keV) in the range 10$^{-13}-6 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ discovering $\\sim$2000 new X-ray sources, $\\sim 80$\\% of which are expected to be active galactic nuclei (AGN). The ChaMP has two parts, the extragalactic survey (ChaMP) and the galactic plane survey (ChaMPlane). Over five years of Chandra operations, the ChaMP will provide both a major resource for Chandra observers and a key research tool for the study of the cosmic X-ray background (CXRB) and the individual source populations which comprise it.  ChaMP promises profoundly new science return on a number of key questions at the current frontier of many areas of astronomy including (1) locating and studying high redshift clusters and so constraining cosmological parameters, (2) defining the true population of AGN, including those that are absorbed, and so constraining the accretion history of the universe, (3) filling in the gap in the luminosity/redshift plane between Chandra deep and previous surveys in studying the CXRB, (4) studying coronal emission from late-type stars as their cores become fully convective and (5) search for cataclysmic variables (CVs) and quiescent Low-Mass X-ray Binaries (qLXMBs) to measure their luminosity functions.  In this paper we summarize the status, predictions and initial results from the X-ray analysis and optical imaging. ",
        "introduction": "The large fields-of-view (FoVs) typical of X-ray imaging instruments have long been exploited for serendipitous surveys, resulting in far-reaching and fundamental advances in our knowledge of the X-ray universe and indeed the universe as a whole ({\\it e.g.}, the Einstein Medium Sensitivity Survey (EMSS, Stocke et al. 1991, Maccacaro et al. 1982); the Cambridge-Cambridge ROSAT Serendipitous Survey (CCRSS, Boyle, Wilkes \\& Elvis 1997); ROSAT International X-ray/Optical Survey (RIXOS, Page et al. 1996)). The next decade promises to be rich for X-ray astronomy as two major observatories, Chandra and XMM, open new vistas with their deeper flux limits and superior spatial and spectral resolution.  \\begin{figure}[h] \\plotfiddle{compare.eps}{2.5in}{-90}{33}{33}{-200}{210} \\vskip -2.0in \\hskip 3.0in \\begin{minipage}{2.5in} \\caption{\\it Predicted source counts per flux bin for ChaMP Cycle 1.\\\\} \\end{minipage} \\vskip 0.85in \\end{figure}  To take full advantage of the rich serendipitous dataset afforded by Chandra's large FoVs (ACIS-I: 16$'\\times 16'$) and to maximize the science per ksec of observing time, we have organized an extragalactic serendipitous X-ray survey, the ChaMP and a similar Galactic plane survey (ChaMPlane). We estimate a ChaMP X-ray sample of $\\sim\\,2000$ sources in the $\\sim\\,5$ sq. degs. of ChaMP fields in Cycle-1. $\\sim 80$\\% of these are likely to be AGN of various types, including those with low luminosity and/or substantial low energy X-ray absorption expected based on the spectrum of the CXRB and already being found in early results from Chandra surveys (Hornschemeier et al. 2000, Mushotzky et al. 2000, Barger et al. 2000, Giacconi et al. 2000). The remaining sources will include clusters of galaxies, stars, X-ray binaries, and supernova remnants. Unlike targeted, small area, deep Chandra surveys, which reach fainter flux levels, the ChaMP provides the wide area essential for finding rare, bright and/or unexpected sources, generating statistically meaningful samples of rare source types such as BL Lac objects, quiescent X-ray binaries and high-redshift clusters. As such the ChaMP complements existing deep Chandra surveys (Figure~1). ",
        "conclusions": ""
    },
    "0011/astro-ph0011141_arXiv.txt": {
        "abstract": "We present a combined maximum-entropy method (MEM) and Mexican Hat wavelet (MHW) analysis in order to recover the different components of the microwave sky. We apply this technique to simulated observations by the ESA Planck satellite in small patches of the sky. In particular, the introduction of the MHW allows one to detect and subtract the brightest point sources present in the input data and therefore to improve the reconstructions of the CMB and foreground components achieved by MEM on its own. In addition, a point source catalogue at each Planck frequency is produced, which is more complete and accurate than those obtained by each technique independently. ",
        "introduction": "Cosmic Microwave Background (CMB) observations carry a wealth of information about the Universe. Indeed, an accurate knowledge of the CMB anisotropies can place tight constraints on fundamental parameters as well as to discriminate between competing theories of structure formation. Future CMB experiments such as the NASA MAP satellite and the Planck mission from ESA, will provide with multifrequency data at high resolution and sensitivity. However, these data contain not only the cosmological signal but also Galactic foregrounds, extragalactic point sources, thermal and kinetic Sunyaev-Zelodvich (SZ) emission from cluster of galaxies and instrumental noise. Therefore our capacity to recover all the valuable information encoded in the CMB will critically depend on our ability to denoise and separate the cosmological signal from the rest of components of the microwave sky.  To perform such a separation,~\\cite{h98} has developed a Fourier MEM algorithm. This technique is particularly successful at using multifrequency data to identify foreground emission from physical components whose spectral signatures are (reasonably) well-known. Therefore, the most problematic foreground to remove is that due to extragalactic point sources, since each source has a unique frequency spectrum and, moreover, is notoriously difficult to predict. To adress this problem,~\\cite{h99} extended the MEM approach to deal with point sources as an extra `noise' contribution. A different approach was developed by~\\cite{sanz} who showed that the Mexican Hat wavelet (MHW) is in fact the optimal pseudo-filter for detecting point sources under reasonable conditions. The application of this wavelet to realistic simulations was presented in~\\cite{c00} and extended in~\\cite{v00}.  The aim of this work is to show that the MEM and MHW techniques are actually complementary and can be combined to improve the accuracy of the separation of diffuse foregrounds from the CMB and increase the number of point sources that are identified and successfully subtracted. The joint analysis has been performed on simulated observations by the Planck satellite but it could be straightforwardly applied to other multifrequency CMB experiments such as the forthcoming NASA MAP satellite or the recently performed Boomerang and MAXIMA experiments. ",
        "conclusions": "We have presented a combined analysis of the maximum-entropy method and the Mexican Hat wavelet to separate and reconstruct the physical components of the microwave sky from multifrequency observations of the CMB that contain point sources. We have applied this technique to simulated data of the Planck satellite pointing out the improvements achieved due to the complementary nature of both approaches. Bright point sources are identified and subtracted by the MHW whereas MEM is able to deal with fainter point sources as a generalised noise. As a result, the reconstructions of the CMB, SZ effects and Galactic foregrounds are improved and mostly free of contaminating point sources. Moreover, using the joint analysis more complete and accurate point source catalogues are produced at each observing frequency as compared to those obtained by each of the techniques independently."
    },
    "0011/astro-ph0011188_arXiv.txt": {
        "abstract": "Interest in non-standard recombination scenarios has been spurred by recent cosmic microwave background (CMB) results from BOOMERANG and MAXIMA, which show an unexpectedly low second acoustic peak, resulting in a best-fit baryon density that is 50\\% larger than the prediction of big-bang nucleosynthesis (BBN). This apparent discrepancy can be avoided if the universe has a non-standard ionization history in which the recombination of hydrogen is significantly delayed relative to the standard model. While future CMB observations may eliminate this discrepancy, it is useful to develop a general framework for analyzing non-standard ionization histories.  We develop such a framework, examining non-standard models in which the hydrogen binding energy $E_b$ and the overall expression for the time rate of change of the ionized fraction of electrons are multiplied by arbitrary factors. This set of models includes a number of previously-proposed models as special cases. We find a wide range of models with delayed recombination that are able to fit the CMB data with a baryon density in accordance with BBN, but there are even allowed models with {\\it earlier} recombination than in the standard model. A generic prediction of these models is that the third acoustic CMB peak should be very low relative to what is found in the standard model. This is the case even for the models with earlier recombination than in the standard model, because here the third peak is lowered by an increased diffusion damping at recombination relative to the standard model. Interestingly, the specific height of the third peak depends sensitively on the model parameters, so that future CMB measurements will be able to distinguish between different non-standard recombination scenarios. ",
        "introduction": "In the past year, observations of the cosmic microwave background (CMB) fluctuations by the BOOMERANG \\cite{boom} and MAXIMA \\cite{max} experiments have produced data of unprecedented precision on CMB fluctuations at small angular scales.  While generally confirming the adiabatic, flat ($\\Omega \\approx 1$) model predicted by inflation, these observations have several puzzling features.  In particular, the position of the first acoustic peak is at a slightly larger angular scale than is predicted in the flat model, and the amplitude of the second peak is unexpectedly low (see, e.g., Ref. \\cite{white}).  If these results are fit using the standard set of cosmological parameters, the result is a CMB prediction for the baryon density of $\\Omega_b h^2 \\sim 0.03$ \\cite{Lange} - \\cite{Tegmark}. In contrast, the prediction for $\\Omega_b h^2$ from Big-Bang nucleosynthesis is $\\Omega_b h^2 \\sim 0.02$ \\cite{OSW,BNT}.  This apparent discrepancy could easily vanish in the light of future CMB measurements.  For the time being, however, it has led to a great deal of interest in models with non-standard ionization histories, since one way to explain the CMB observations and preserve agreement with the BBN baryon density is to postulate that the epoch of recombination was delayed to a lower redshift than in the standard model.  Peebles, Seager, and Hu suggested that this could occur if there were sources of Ly $\\alpha$ photons present at the epoch of recombination \\cite{peebles}. A more speculative mechanism is a time-variation in the fine-structure constant, $\\alpha$ \\cite{hannestad}-\\cite{landau}.  The authors of Ref. \\cite{landau} used such a time-variation (along with changes in the cosmological parameters) as an example of how one might model non-standard recombination in general.  Because there are a variety of models with non-standard ionization histories, we feel that is it worthwhile to try to develop a general framework for analyzing such models. Ideally, one would like to investigate the consequences for the CMB of an arbitary $x_e(z)$, the ionization fraction as a function of redshift. It is obviously impractical to investigate arbitrary functional forms for $x_e$.  Instead, we have attempted to parametrize deviations of $x_e(z)$ from the standard ionization history in terms of a small number of physically-motivated parameters.  In particular, we multiply the overall ionization/recombination rates by a free parameter $a$, and the binding energies by a parameter $b$. We then calculate the predictions for the CMB fluctuation spectrum as a function of these parameters and compare with observations. Our parametrization is discussed in the next section, along with the results of modifying the ionization history. Our conclusions are given in Sec. 3.  Although we have compared our results with the recent BOOMERANG and MAXIMA experiments, our work provides a general framework for discussing non-standard recombination scenarios, and it can be applied to future observations as well. ",
        "conclusions": "Our investigation of the effect of a non-standard ionization history on the CMB using a change in the overall ionization/recombination rates (through the parameter $a$) and in the binding energies (through the parameter $b$) appears to provide a very general framework for studying such variations in the standard model.  In particular, it seems possible to model arbitrary changes in both the epoch of last scattering and in the width of the last-scattering surface through such variations, with the exception of ionization histories having a narrow, high-redshift surface of last scattering, for which the ionization fraction decreases faster than in the equilibrium case.  Our results indicate that the BOOMERANG and MAXIMA results with $\\Omega_b h^2 = 0.019$ (from BBN) can be well-fit over a broad range of choices for $a$ and $b$. However, the lower bounds on $a$ and $b$ are quite robust (i.e., nearly independent of variation in the other parameter). We find that  $b >0.5$ and $\\log(a) > -1.6$ at the 95\\% confidence level. If we restrict $a$ to its standard model value, changing only the binding energies, the allowed (95\\%) region is given by $0.5 < b < 1.1$. Similarly, if we restrict the binding energies to be unchanged ($b = 1$) and change only the overall rates, then $a$ is constrained to lie in the range $-1.6 < \\log(a) < 0.4$.  Our best-fit model for $\\Omega_b h^2 = 0.019$ has a $\\chi^2$ per degree of freedom of 0.50.  In comparison, the best fit for standard reionization ($a=b=1$) with $\\Omega_b h^2 = 0.030$ gives $\\chi^2$ per degree of freedom of 0.82, so both models are good fits.  (These numbers include a 10\\% calibration error for BOOMERANG and 4\\% for MAXIMA).  One robust prediction of all of these models is a decrease in the amplitude of the third peak, relative to its height in the standard model.  In Fig. 9, we show the ratio $X$ of the height of the first peak to the height of the third peak as a function of $a$ and $b$. The standard model ($a=b=1$, $\\Omega_b h^2 =0.019$) value, $X = 2.65$, is shown as a dashed curve. The entire $1\\sigma$ region gives values for $X$ larger than in the standard model.  In comparison, the best fit model with high baryon density ($\\Omega_b h^2 = 0.030$) and a standard recombination history produces a much higher third peak ($X = 1.82$) than in the standard model.  Our results are more general than previous comments that models with late recombination can be distinguished from high $\\Omega_b h^2$ models by the amplitude of the third peak.  Our allowed region includes models with recombination at slightly {\\it higher} redshift than in the standard model (these models lie on the extreme left-hand side of Fig. 9).  For these models, the third peak is still reduced in amplitude, but this reduction is due to diffusion damping from an increase in the width of the surface of last scattering, rather than from a decrease in the redshift of last scattering as in Refs. \\cite{peebles} - \\cite{landau}.  Our results seem to imply that any modification in the recombination history which fits the BOOMERANG and MAXIMA observations will result generically in a decrease in the height of the third peak.  However, we also find that the actual height of the third peak is a function of $a$ and $b$.  From Fig. 9 it can be seen that this ratio increases as we move through the allowed region from small $a$, large $b$ to large $a$, small $b$.  So if the universe did have a non-standard ionization history, the amplitude of the third peak as shown in Fig. 9 should allow us to determine precisely the nature of the deviation of $x_e(z)$ from its standard evolution.   \\vskip 0.2in \\noindent We thank U. Seljak and M. Zaldariagga for the use of CMBFAST \\cite{CMBFAST}. R.J.S. was supported in part by the DOE (DE-FG02-91ER40690)."
    },
    "0011/astro-ph0011231_arXiv.txt": {
        "abstract": "The current status of the UltraHigh Energy Cosmic Ray (UHE CR) enigma and several proposed solutions involving particle physics beyond the standard model are discussed. Emphasis is given to top--down models, and as a main example, supermassive dark matter as galactic source for UHE CR and the status of its experimental signatures (galactic anisotropy, chemical composition and clustering) is reviewed. Then different approaches to calculate fragmentation spectra of supermassive particles are discussed. Finally, it is argued that UHE neutrinos cannot be -- neither directly or indirectly -- responsible for the observed vertical air showers. ",
        "introduction": "Cosmic Rays~\\cite{CR} (CR) are observed in a wide energy range, starting from subGeV energies up to $3\\times 10^{20}$~eV. Apart from the highest energies, these particles are accelerated in our Galaxy, most probably by shocks produced by SNII explosions. Since the galactic magnetic field cannot confine and isotropize particles with energy higher than $\\sim Z\\times 10^{19}$~eV, it is natural to think that the UltraHigh Energy (UHE) component has an extragalactic origin. Moreover, the acceleration of protons or nuclei up to 2--3$\\times 10^{20}$~eV is difficult to explain within the known astrophysical galactic sources.  The most prominent signature of extragalactic UHECR is the so called Greisen--Zatsepin--Kuzmin (GZK) cutoff~\\cite{GZK}: the energy losses of protons sharply increase at $E_{\\rm GZK} \\ap 6\\times 10^{19}$~eV, owing to scattering on photons from the microwave background, $p+\\gamma_{3K}\\to\\Delta^+ \\to N+\\pi$, reducing their mean free path length to less than $50$~Mpc or so. Nuclei exhibit a similar cutoff at the same energy, while photons have no sharp cutoff but an even shorter free mean path, because of pair-production on the radio background. Thus, the UHECR spectrum should dramatically steepen above $E_{\\rm GZK}$ for {\\em any\\/} homogeneous distribution of proton or nuclei sources. However, the observed spectrum extends up to $3\\times 10^{20}$~eV and this maximum seems to be caused by the limited exposure. Both spectra, the expected and the observed one, are shown in Fig.~1, taken from Ref.~\\cite{agasa00}.   There is another argument that disfavours the standard astrophysical sources: At energies $E\\sim 10^{20}$~eV, the arrival direction of the primaries which is known within several degrees should point towards their site of origin. But no source of UHECR such us active galactic nuclei (AGN) has been found within $50$~Mpc in the direction of these events. In Fig.~\\ref{cluster}, the arrival directions of 92 events above $4\\times 10^{19}$~eV are shown in galactic coordinates~\\cite{Uc00}. The events are scattered isotropically on larger scales, and no significant enhancement towards the galactic or supergalactic plane is found. Intriguingly, about 20\\% of the events are clustered in angular doublets or even triplets; both triplets are found near the supergalactic plane.  \\begin{figure} \\begin{center} \\epsfig{file=agasa_spec00.eps, height=7cm,width=10cm} \\caption{Energy spectrum as observed by AGASA. Error bars are 68\\% C.L., numbers are the number of events per energy bin. The dashed line is the spectrum expected from uniformly distributed astrophysical sources (from Ref.~[3]). \\label{agasa}} \\end{center} \\end{figure}  An elegant solution to some of the problems described above are top--down models: in contrast to the standard sources, the primaries are not accelerated but are the fragmentation products of some decaying superheavy particle $X$. For $X$-particles with mass $m_X\\gsim 10^{12}$~GeV, the acceleration problem is solved trivially. Moreover, these sources also evade detection by normal astronomical methods.   This review is organised as follows: In Sec.~II top-down models and in Sec.~III their signatures are presented. We then discuss how the spectrum of hadrons produced in the decay of supermassive particles can be calculated. Finally, we argue  in Sec.~5 that UHE neutrinos cannot be -- neither directly or indirectly via resonant annihilations on relic neutrinos -- responsible for the observed UHECR events.  For a review of astrophysical sources, we refer the interested reader to Ref.~\\cite{astro}. The suggestion that the violation of Lorentz invariance induced by quantum gravity effects avoids the GZK cutoff is covered in the recent review~\\cite{El00}.   \\begin{figure} \\begin{center} \\epsfig{file=cluster.eps, height=7cm,width=10cm} \\caption{\\label{cluster} Map of 92 events above $4\\times 10^{19}$~eV in galactic coordinates (from Ref.~[4]). The supergalactic plane is shown by the dotted line, the hatched region is unobservable.} \\end{center} \\end{figure} ",
        "conclusions": "Superheavy, metastable relic particles is the most promising source in the framework of top--down models. Its cleanest signature is the galactic anisotropy which should be easily detectable by a detector in the southern hemisphere as AUGER. Experimentally more difficult signatures are the high photon/proton ratio and the detection of the predicted large UHE neutrino fluxes via horizontal EAS. All top--down models predict also an appreciable LSP flux if R-parity is conserved.  Also in theories with large extra dimensions, neutrinos behave as deeply penetrating particles and are therefore not responsible for the observed UHECR events."
    },
    "0011/astro-ph0011007_arXiv.txt": {
        "abstract": "Here we introduce the Delaunay Density Estimator Method. Its purpose is rendering a fully volume-covering reconstruction of a density field from a set of discrete data points sampling this field. Reconstructing density or intensity fields from a set of irregularly sampled data is a recurring key issue in operations on astronomical data sets, both in an observational context as well as in the context of numerical simulations. Our technique is based upon the stochastic geometric concept of the Delaunay tessellation generated by the point set. We shortly describe the method, and illustrate its virtues by means of an application to an N-body simulation of cosmic structure formation. The presented technique is a fully adaptive method: automatically it probes high density regions at maximum possible resolution, while low density regions are recovered as moderately varying regions devoid of the often irritating shot-noise effects.  Of equal importance is its capability to sharply and undilutedly recover anisotropic density features like filaments and walls.  The prominence of such features at a range of resolution levels within a hierarchical clustering scenario as the example of the standard CDM scenario is shown to be impressively recovered by our scheme. ",
        "introduction": "Astronomical observations, physical experiments as well as computer simulations often involve discrete data sets supposed to represent a fair sample of an underlying smooth and continuous field. Conventional methods are usually plagued by one or more artefacts. Firstly, they often involve estimates at a restricted and discrete set of locations -- usually defined by a grid -- instead of a full volume-covering field reconstruction. A problem of a more fundamental nature is that the resulting estimates are implicitly mass-weighted averages, whose comparison with often volume-weighted analytical quantities is far from trivial. For most practical purposes, the disadvantage of almost all conventional methods is their insensitivity and inflexibility to the sampling point process. This leads to a far from optimal performance in both high density and low density regions, which often is dealt with by rather artificial and ad hoc means.  In particular in situations of highly non-uniform distributions conventional methods tend to conceal various interesting and relevant aspects present in the data. The cosmic matter distribution exhibits conspicuous features like filaments and walls, extended along one or two directions while compact in the other(s). In addition, the density fields display structure of varying contrasts over a large range of scales. Ideally sampled by the data points, appropriate field reconstructions should be set solely and automatically by the point distribution itself. The commonly used methods, involving artificial filtering through for instance grid size or other smoothing kernels (e.g. Gaussian filter) often fail to achieve an optimal result.  \\begin{figure*} \\begin{picture}(100,150)(0,10) \\includegraphics[width=12cm,keepaspectratio]{areapanelbw.eps} \\end{picture} \\hfill \\parbox[b]{5.5cm}{ \\caption{\\small A set of 20 points with their Voronoi (left frame: solid lines) and Delaunay (right frame: solid lines) tesselations. Left frame: the shaded region indicates the Voronoi cell corresponding to the point located just below the center. Right frame: the shaded region is the ``contiguous Voronoi cell'' of the same point as in the lefthand frame. }\\vspace*{-.3cm} } \\end{figure*}  Here we describe and propagate a new fully self-adaptive method based on the Delaunay triangulation of the given point process. After a short description of the fundamentals of our tessellation procedure, we show its convincing performance on the result of an N-body simulation of structure formation, whose particle distribution is supposed to reflect the underlying cosmic density field. A detailed specification of the method, together with an extensive quantitative and statistical evaluation of its performance will be presented in a forthcoming publication \\cite{schetwey20}. ",
        "conclusions": ""
    },
    "0011/astro-ph0011013_arXiv.txt": {
        "abstract": "We present a semi-empirical model for the evolving far-infrared to ultraviolet diffuse background produced by stars in galaxies. The model is designed to reproduce the results of deep galaxy surveys, and therefore may be considered as a cosmology-independent lower limit to the extragalactic background light. Using this model and recent HEGRA data, we infer the intrinsic spectrum at multi-TeV gamma-ray energies for Mkn~501 and find that it is consistent with a power law of spectral index $2.49\\pm0.04$.  In turn, this finding renders it rather unlikely that the present-day infrared background has an intensity as high as claimed by Finkbeiner et al.~\\cite{lit:finkbeiner}. Future 10~GeV to TeV observations could be used to either constrain the ultraviolet-to-infrared background model at high redshifts or cosmological parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011539_arXiv.txt": {
        "abstract": "#1{\\vskip 7mm \\begin{center}{\\large Abstract}\\par \\smallskip \\begin{minipage}[c]{12cm} \\small #1 \\end{minipage} \\end{center} } \\def\\title#1{\\begin{center}{\\Large\\bf #1}\\end{center}} \\def\\author#1{\\vskip 5mm \\begin{center}{#1}\\end{center}} \\def\\address#1{\\begin{center}{\\it #1}\\end{center}} \\newcommand{\\del}{\\partial} \\newcommand{\\m}{\\mathbf} \\newcommand{\\x}{{\\mathbf{x}}} \\newcommand{\\y}{{\\mathbf{y}}} \\newcommand{\\z}{{\\mathbf{z}}} \\newcommand{\\n}{{\\mathbf{n}}} \\newcommand{\\k}{{\\mathbf{k}}} \\newcommand{\\U}{\\underline}  \\newcommand{\\ti}{\\textit} \\newcommand{\\f}{\\frac} \\newcommand{\\T}{\\tilde} \\newcommand{\\N}{\\nonumber} \\newcommand{\\bb}{\\bibitem} \\newcommand{\\BF}{\\begin{figure}\\begin{center}} \\newcommand{\\EF}{\\end{center}\\end{figure}} \\newcommand{\\BE}{\\begin{equation}} \\newcommand{\\EE}{\\end{equation}} \\newcommand{\\BEA}{\\begin{eqnarray}} \\newcommand{\\EEA}{\\end{eqnarray}} \\makeatletter \\@ifundefined{lesssim}{\\def\\lesssim{\\mathrel{\\mathpalette\\vereq<}}}{} \\@ifundefined{gtrsim}{\\def\\gtrsim{\\mathrel{\\mathpalette\\vereq>}}}{} \\def\\vereq#1#2{\\lower3pt\\vbox{\\baselineskip1.5pt \\lineskip1.5pt \\ialign{$\\m@th#1\\hfill##\\hfil$\\crcr#2\\crcr\\sim\\crcr}}} \\makeatother  \\begin{document}  \\title{ Cosmic Microwave Background in Closed Multiply Connected Universes } \\author{ Kaiki Taro Inoue,\\footnote{E-mail:tinoue@yukawa.kyoto-u.ac.jp} } \\address{ Yukawa Institute for Theoretical Physics, Kyoto University, \\\\ Kyoto 606--8502, Japan } \\abstract{We have investigated the cosmic microwave background (CMB) anisotropy in closed multiply connected universes (flat and hyperbolic) with low matter density. We show that the COBE constraints on these low matter density models with non-trivial topology are less stringent since a large amount of CMB anisotropy on large angular scales can be produced due to the decay of the gravitational potential at late time. } ",
        "introduction": "For a long time, cosmologists have assumed the simply connectivity of the spatial hypersurface of the universe. If it is the case, the topology of closed 3-spaces is limited to that of a 3-sphere if Poincar\\'{e}'s conjecture is correct. However, there is no particular reason for assuming the simply connectivity since the Einstein equations do not specify the boundary conditions. If we allow the spatial hypersurface being multiply connected then the spatial geometry of closed models can be flat or hyperbolic as well. We should be able to observe the imprint of the ``finiteness'' of the spatial geometry if it is multiply connected on scales of the order of the particle horizon or less, in other words, if we live in a ``small universe''. \\\\ \\indent However, it has been claimed by several authors that the ``small universes'' have already been ruled out observationally. For a flat 3-torus model without the cosmological constant, the COBE data constrains the topological identification scale $L$ (the minimum length of periodic geodesics) to be $L\\ge 0.4\\times 2 R_\\ast$ where $R_\\ast$ is the comoving radius of the last scattering surface \\cite{Sokolov,Starobinsky,Stevens,Oliveira1}. The suppression of the fluctuations on scales beyond $L$ leads to a decrease of the angular power spectrum $C_l$ of the CMB temperature fluctuations on large angular scales. \\\\ \\indent In contrast, for low matter density models, the constraint could be considerably less stringent since a bulk of large-angle CMB fluctuations can be produced by the so-called (late) integrated Sachs-Wolfe effect \\cite{Cornish2} which is the gravitational blueshift effect of the free streaming photons caused by the decay of the gravitational potential\\cite{HSS}. If the background geometry is either flat or hyperbolic, then the angular size of a fluctuation becomes small as it approaches to the observation point. Although fluctuations beyond the size of the fundamental domain are suppressed, no significant suppression in large-angle power occurs if they are produced at place well after the last scattering. In other words, large-angle fluctuations can be generated when the fluctuations enter the topological identification scale at late time. ",
        "conclusions": "We have investigated the CMB anisotropy in closed multiply connected universes (flat and hyperbolic) with low matter density. We have seen that the COBE constraints for these models are less stringent compared with the simplest ``standard'' 3-torus model ($\\Omega_m\\!=\\!1.0$). On the other hand recent observations of distant supernova Ia \\cite{Perlmutter,Riess} and of the CMB on small angular scales \\cite{Boomerang,MAXIMA} imply that our universe is nearly flat with the cosmological constant (or ``quintessence'', X-matter etc.) which dominates the present universe. It seems that the closed hyperbolic models are ruled out but it is still not conclusive. If one includes the cosmological constant for a fixed curvature radius, the comoving radius of the last scattering surface in unit of curvature radius becomes large. Therefore the observable effects of the non-trivial topology become much prominent. For instance, the number $N$ of copies of the fundamental domains inside the observable region at present is approximately 28 for the Weeks model with $\\Omega_\\Lambda\\!=\\!0.6$ and $\\Omega_m\\!=0.2\\!$ whereas $N\\!=\\!4$ if $\\Omega_\\Lambda\\!=\\!0$ and $\\Omega_m\\!=0.8\\!$. If one allows hyperbolic orbifold models then the number can be increased as much as $N=52$ for $\\Omega_\\Lambda\\!=\\!0.75$ and $\\Omega_m\\!=\\!0.2$ since the volume of the smallest orbifold (arithmetic) is very small (=0.039). At least at the classical level it seems that there is no reason to exclude any orbifold models although they have singular points where the curvature diverges."
    },
    "0011/astro-ph0011155_arXiv.txt": {
        "abstract": "\\zwb\\ (UGC~9562) was classified as a ``probable'' polar-ring galaxy in the Polar-Ring Catalog (\\cite{PRC}), based upon its optical appearance. We present 21\\,cm and optical observations of this galaxy and its companion, the blue star-forming dwarf \\zwa. Our 21\\,cm observations show that % $5 \\times 10^8$ solar masses of \\HI\\ is present in a polar ring orbiting \\zwb, and show a spatially and kinematically contiguous streamer with $2.5 \\times 10^8$ solar masses of \\HI\\ gas between the two galaxies.  This gaseous bridge, plus our observations of \\Halpha\\ line emission in the polar ring and in \\zwb, are strong evidence for an ongoing interaction between the two galaxies.  However, the configuration of the streamer suggests that the polar ring itself may well have predated the current interaction, which then stimulated an outburst of star formation in the ring gas. ",
        "introduction": "\\label{intro}  Polar-ring galaxies (PRG's) are systems with two kinematically distinct components.  The central component (the ``host galaxy'') is usually a small S0 galaxy, or occasionally an elliptical galaxy, and the outer component (the ``polar ring'') is a ring or annulus of material in a nearly circumpolar orbit about the host.  The fact that the angular momentum vectors of the host galaxy and of the polar ring are nearly orthogonal indicates that PRG's are most likely the end products of mergers or accretion events, rather than a single coherent formation process.  The Polar-Ring Catalog (\\cite{PRC}, hereafter PRC) classifies 6 galaxies as kinematically-confirmed polar rings (category ``A'' in the PRC), and 27 galaxies as ``probable'' polar rings based upon their optical appearance (category ``B'' in the PRC).  The catalog is frozen, but some category ``B'' galaxies have since been confirmed since its publication (see \\cite{SC00}).  The origin of the material in polar rings is something of a mystery, because the gas masses involved are so large.  Polar rings are gas-rich, with typically a few billion solar masses of \\HI\\ gas (\\cite{vGSK87}; \\cite{RSS94}; \\cite{CS96}; \\cite{A97}), and often significant molecular gas and dust as well (see \\cite{GSS97}).  The identified polar rings have enough stars that we can see the rings at optical wavelengths, and stellar rings lie within the rings of neutral gas (\\eg\\, \\cite{vGSK87}; \\cite{A97}).  Because accreted gas can dissipate energy while the stars cannot, gas and stars will tend to settle into very different final configurations. The stars now in the polar ring must have been formed from the gas after it had settled into the thin ring structure (\\eg, \\cite{DC94}).  The host galaxies are small early-type galaxies, typically with L$_B$~$<$~10$^9$~\\Msun.  There is no evidence that the host galaxies themselves contain substantially more cool gas and dust than other early-type galaxies without polar rings (see \\cite{RSS94}).  Their radio continuum emission is not significantly stronger than early-type galaxies without polar rings (\\cite{CS94}; \\cite{Cetal00}). This lack of radio continuum emission indicates that any strong starburst activity which may have been induced by the event that formed the ring has since ceased.  Even making no allowance for the mass of stars which have formed in the polar rings, these galaxies have captured the equivalent of the entire gas content of a typical spiral galaxy (\\eg, \\cite{RH94}, hereafter RH94).  It is not clear that a large spiral galaxy could accrete onto a small S0 galaxy without leaving behind a very disorganized system.  If the spiral companion was destroyed, its stars would not dissipate kinetic energy to form a ring as the gas has done, but should be visible as faint shells or diffuse light around the merged system.  Although some PRG's are known to have shells, many PRG's do not have visible shells (PRC).  If the ring material was taken from the outer parts of a spiral galaxy, then we should be able to see the remains of this companion, but many polar rings have no nearby companions.  If the polar ring was formed by accretion of a gas-rich dwarf galaxy, there might be no conspicuous stellar shells, and no visible remnant of the encounter.  But dwarf irregular galaxies at the present day contain too little neutral gas to make polar rings; the upper end of the \\HI\\ mass function for present-day dwarf galaxies is only a few times 10$^8$~\\Msun\\ (\\eg\\ \\cite{BR93}; \\cite{Br97}; \\cite{MvDG98a}).  Searches for clouds of primordial intergalactic \\HI\\ indicate that there is very little gas in this form at the current epoch (\\eg\\ \\cite{Br93}; \\cite{Tetal96}; \\cite{Zetal97}).  If PRG's are stable structures, most of them may have formed early in the history of the universe, when large clouds of intergalactic gas should have been more common.  In any case, it is highly probable that the accretion of material onto many host galaxies is not complete.  The gas of the polar ring has come from an orbit of high angular momentum with a large impact parameter, and it is likely that streamers of gas would remain around most newly-formed polar rings well after the original encounter (\\cite{HM95}).  The study of probable polar-ring systems which have nearby companions are therefore very likely to reveal gaseous tails or bridges which may help us better understand the interactions which form polar rings.  \\zwb\\ (UGC~9562) is classified as ``B-17'', a ``probable'' polar ring system, in the PRC.  An optical image of this galaxy and its blue dwarf companion, \\zwa\\, is shown in Figure~\\ref{rcont}; the polar ring runs NE to SW.  Optical spectroscopy in the \\Halpha\\ emission line of \\zwb\\ by Reshetnikov \\&\\ Combes (1994) revealed two gaseous components with orbits nearly perpendicular to each other --- one with a velocity gradient along the apparent major axis of the polar ring, and another with a velocity gradient along the apparent major axis of the host galaxy.  \\zwb\\ is thus a kinematically-confirmed polar-ring galaxy. Earlier observations of this system by Balkowski \\etal\\ (1978; hereafter BCW78) in the 21\\,cm line with the Westerbork Synthesis Radio Telescope (WSRT) also showed rotation of gas along the polar ring, and in addition detected a cloud of \\HI\\ gas between the two galaxies.  No optical counterpart brighter than 25.5~mag\\,arcsec$^{-2}$ to this cloud was detected in {\\em B} band.  The \\zwab\\ system is thus one of the best known candidates for a polar ring in the process of formation by accretion of gas from a nearby companion.  We have obtained observations in the 21\\,cm line of \\HI, using the Very Large Array (VLA\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}) aperture synthesis telescope.  These observations are more sensitive than the observations by BCW78, and have approximately twice the spatial and velocity resolution.  We have also obtained \\Halpha\\ and broad-band optical images of this system, to search for evidence of ongoing star formation and for an optical counterpart of the gas cloud between \\zwa\\ and \\zwb.  The \\HI\\ distribution and kinematics of the \\zwab\\ system are presented in Section~\\ref{21cm}, and the optical data in Section~\\ref{optical}.  In Section~\\ref{discussion}, we discuss the morphology of the stars and gas in this system, and investigate the possibility that the polar ring around \\zwb\\ is the product of an ongoing accretion event from gas originating in \\zwa.  We take \\Ho~=~75~km~s$^{-1}$~Mpc$^{-1}$, implying a distance of 18.1~Mpc for the system. ",
        "conclusions": "\\label{discussion}   We observe \\HI\\ gas in regular rotation about the minor axis of the optical polar ring in \\zwb, and none that shares the rotation of the S0 disk.  \\zwb\\ is thus a confirmed polar-ring galaxy.  The amount of \\HI\\ is fairly modest, compared with other polar rings (\\cite{RSS94}), but is much greater than normal for an S0 galaxy, and is more typical of a late-type spiral or large irregular galaxy (\\cite{RH94}).  As with other polar rings, the extent of the \\HI\\ gas is greater than that of the optical ring.  The rotation curve does not drop, even far outside the optical radius of the host galaxy.  As in other polar-ring galaxies, it appears that most of the mass in this galaxy is dark (\\eg, \\cite{SWR83}; \\cite{WMS87}; \\cite{SRJF94}; \\cite{A97}).   The polar ring of \\zwb\\ is asymmetric, at both optical and radio wavelengths.  In addition, observations at \\Halpha\\ show clumpy emission in the ring.  This emission is more pronounced on the northern side of the galaxy, the same side on which faint, steep-spectrum radio continuum emission is observed.  The presence of both \\Halpha\\ emission and radio continuum emission along the ring, and its blue color, are strong arguments for ongoing star formation. The asymmetry of the ring material and the evidence of star formation in the ring make it likely that new material has recently been added to the polar ring, or that pre-existing ring material has recently been disturbed by a new interaction.   At our assumed distance of 18.1~Mpc, the projected separation of 4.4\\arcmin\\ between these two galaxies corresponds to a separation of $\\simeq$~23~~kpc.  At a relative speed of 100~\\kms, it would have taken these galaxies at least 230~Myr to separate after a close interaction.  The orbital timescale for gas in the outer regions of \\zwb\\ is about 200~Myr, implying that the most recent interaction between these two galaxies may have taken place within a few orbital timescales.  Was the polar ring of \\zwb\\ formed by accretion from the companion galaxy \\zwa?  The two galaxies are smoothly connected in both position and velocity by a gaseous bridge or streamer, in which we observe no stellar component.  Assuming that the bridge is the remnant of the interaction that created the polar ring, so that all of the gas from the polar ring and streamer originally belonged to \\zwa, would yield $\\rm{M_{HI}/L_{B}}~=~0.55$ prior to the interaction --- a high, but not unreasonable, value for an irregular galaxy. Since the dynamical mass of \\zwb\\ is more than twice that of \\zwa, the outer layers of a gas-rich system the size of \\zwa\\ would not have been very tightly bound, and might have been fairly easily detached.   The blue stars and \\Halpha\\ regions in the polar ring are spread along a line, implying that much of the gas has already settled into a disk or ring, that we now see almost edge-on.  It is difficult to explain how gas that was accreted only a few orbital timescales ago could have already settled into a plane.  If there were enough dissipation to cause flattening on this short timescale, we would expect substantial amounts of accreted gas to be falling into the center of the host galaxy, and either forming stars there, or fueling a central black hole.  However, this does not seem to be the case.  Our observations reveal no radio continuum source, and the far-infrared IRAS flux is also low (PRC; \\cite{RSS94}).  The ionized gas of the central galaxy does not share the rotation axis of the ring material (\\cite{RC94}), so it is unlikely to have come from the ring.   The contiguous streamer of gas between \\zwa\\ and \\zwb, as well as the star formation activity in both \\zwa\\ and the polar ring, indicate an ongoing interaction between these two galaxies.  However, the fact that HI gas to the east of the polar ring appears to be a continuation of the streamer to the west, in both position and velocity, implies that the streamer probably lies behind \\zwb\\ or in front of it, rather than merging into the polar ring.  That is not what we would expect if both the polar ring and the streamer originated in the same, recent, accretion event. A further indication that the ring itself has existed for some time comes from the relatively weak \\Halpha\\ emission.  At the present rate of star formation, it would take several gigayears to build up the stars that we now observe in the ring, while the blue color indicates that starbirth was no more rapid in the past than it is now. We suggest instead that the polar ring of \\zwb\\ predated the most recent close passage of \\zwa, but was disturbed by it.  Just as in the disks of spiral galaxies (Larson \\& Tinsley 1978, Mihos \\& Hernquist 1996, Barton, Geller \\& Kenyon 2000), tidal torques have `rejuvenated' the ring, causing the vigorous star formation that we now observe."
    },
    "0011/astro-ph0011319_arXiv.txt": {
        "abstract": "We present a quantitative model for the infrared emission from dust in the diffuse interstellar medium. The model consists of a mixture of amorphous silicate grains and carbonaceous grains, each with a wide size distribution ranging from molecules containing tens of atoms to large grains $\\gtsim 1$ micron in diameter. We assume that the carbonaceous grains have polycyclic aromatic hydrocarbon (PAH)-like properties at very small sizes, and graphitic properties for radii $a \\gtrsim 50$\\AA. On the basis of recent laboratory studies and guided by astronomical observations, we propose ``astronomical'' absorption cross sections for use in modeling neutral and ionized PAHs from the far ultraviolet to the far infrared. We also propose modifications to the far-infrared emissivity of ``astronomical silicate''.  We calculate energy distribution functions for small grains undergoing ``temperature spikes'' caused by stochastic absorption of starlight photons, using realistic heat capacities and optical properties. Using a grain size distribution consistent with the observed interstellar extinction, we are able to reproduce the near-IR to submillimeter emission spectrum of the diffuse interstellar medium, including the PAH emission features at 3.3, 6.2, 7.7, 8.6, and 11.3$\\micron$. The model is compared with the observed emission at high Galactic latitudes as well as in the Galactic plane, as measured by the COBE/DIRBE, COBE/FIRAS, IRTS/MIRS, and IRTS/NIRS instruments. The model has $60\\times 10^{-6}$ of C (relative to H) locked up in PAHs, with $45\\times 10^{-6}$ of C in a component peaking at $\\sim 6$\\AA\\ ($N_{\\rm C}\\approx 100$ carbon atoms) to account for the PAH emission features, and with $15\\times 10^{-6}$ of C in a component peaking at $\\sim 50$\\AA\\ to account for the 60$\\mu$m flux. The total infrared emission is in excellent agreement with COBE/DIRBE observations at high galactic latitudes, just as the albedo for our grain model is in accord with observations of the diffuse galactic light. The aromatic absorption features at 3.3$\\mu$m and 6.2$\\mu$m predicted by our dust model are consistent with observations.  We calculate infrared emission spectra for our dust model heated by a range of starlight intensities, from 0.3 to 10$^4$ times the local interstellar radiation field, and we tabulate the intensities integrated over the SIRTF/IRAC and MIPS bands. We also provide dust opacities tabulated from the extreme ultraviolet to submillimeter wavelengths. ",
        "introduction": "}  It has been 70 years since Trumpler's discovery of color excesses provided the first definitive proof of the existence of interstellar dust (Trumpler 1930); however, many aspects of the nature of interstellar dust still remain unclear. Our knowledge of the composition of interstellar grains is based mainly on their spectroscopic absorption/emission features and observed elemental depletions. The most generally accepted view is that interstellar grains consist of amorphous silicates and some form of carbonaceous materials; the former inferred from the 9.7$\\mu$m Si-O stretching mode and 18$\\mu$m O-Si-O bending mode absorption features in interstellar regions as well as the fact that the cosmically abundant heavy elements such as Si, Fe, Mg are highly depleted; the latter mainly inferred from the strongest interstellar absorption feature -- the 2200\\AA\\ hump -- and the fact that silicates alone are not able to provide enough extinction.  The dust sizes, inferred from the wavelength-dependent interstellar extinction and polarization curves, may be separated into two domains -- 1) radii $a > 0.02\\micron$, which includes the ``classical'' grains (with $a \\simgt 0.1\\mu$m), which are primarily responsible for the extinction, polarization and scattering at visible wavelengths; and 2) the ``very small grain'' component (with $a <0.02\\mu$m) which contribute importantly to the extinction in the vacuum-ultraviolet. While the size distribution for the ``classical grain'' component is relatively well constrained by fitting the observed interstellar extinction curve for an assumed dust composition, our knowledge of the size distribution for the ``very small grain'' component is much poorer, due to the fact that, for $\\lambda \\simgt 0.1\\mu$m, these very small grains are in the Rayleigh limit and their extinction cross sections per unit volume are independent of size, so that the observed UV/far-UV extinction curve only constrains the total volume (mass) of this grain component.  Forty-five years ago, Platt (1956) proposed that very small grains or large molecules with radii $\\simlt 10$\\AA\\ may be present in interstellar space. Donn (1968) further proposed that polycyclic aromatic hydrocarbon-like ``Platt particles'', may be responsible for the UV interstellar extinction. However, the inability to determine the detailed optical properties of ``Platt particles'' and the limited observational information on the UV/far-UV interstellar extinction forestalled further development of these ideas (Greenberg 1960).  Since the 1980s, an important new window on the ``very small grain component'' has been opened by infrared (IR) observations. The near-IR continuum emission of reflection nebulae (Sellgren, Werner, \\& Dinerstein 1983) and the 12 and 25$\\mu$m ``cirrus'' emission detected by the {\\it Infrared Astronomical Satellite} (IRAS) (Boulanger \\& P\\'{e}rault 1988) explicitly indicated the presence of a very small interstellar dust component since large grains (with radii $\\sim 0.1\\mu$m) heated by diffuse starlight emit negligibly at such short wavelengths, whereas very small grains (with radii $\\simlt 0.01\\mu$m) can be transiently heated to very high temperatures ($\\simgt 1000$\\ K depending on grain size, composition, and photon energy) as first pointed out by Greenberg (1968). Subsequent measurements by the {\\it Diffuse Infrared Background Experiment} (DIRBE) instrument on the {\\it Cosmic Background Explorer} (COBE) satellite confirmed this and detected additional broadband emission at 3.5 and 4.9$\\mu$m (Arendt et al.\\ 1998).  More recently, spectrometers aboard the {\\it Infrared Telescope in Space} (IRTS) (Onaka et al.\\ 1996; Tanaka et al.\\ 1996) and the {\\it Infrared Space Observatory} (ISO) (Mattila et al.\\ 1996) have shown that the diffuse interstellar medium radiates strongly in emission features at 3.3, 6.2, 7.7, 8.6, and 11.3$\\mu$m. These emission features, first seen in the spectrum of the planetary nebulae NGC 7027 and BD+30$^{\\rm o}$3639 (Gillett, Forrest, \\& Merrill 1973), have been observed in a wide range of astronomical environments including planetary nebulae, protoplanetary nebulae, reflection nebulae, HII regions, circumstellar envelopes, and external galaxies (see Tielens et al.\\ 1999 for a review). Often referred to as ``unidentified infrared'' (UIR) bands, these emission features are now usually attributed to polycyclic aromatic hydrocarbons (PAHs) which are vibrationally excited upon absorption of a single UV/visible photon (L\\'{e}ger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985) although other carriers have also been proposed such as hydrogenated amorphous carbon (Duley \\& Williams 1981; Borghesi, Bussoletti, \\& Colangeli 1987; Jones, Duley, \\& Williams 1990), quenched carbonaceous composite (Sakata et al.\\ 1990), coal (Papoular et al.\\ 1993), fullerenes (Webster 1993), and interstellar nanodiamonds with $sp^3$ surface atoms reconstructed to $sp^2$ hybridization (Jones \\& d'Hendecourt 2000). The emission mechanism proposed for the UIR bands -- UV excitation of gas-phase PAHs followed by internal conversion and IR fluorescence -- is supported by laboratory measurements of the IR {\\it emission} spectra of gas-phase PAH molecules (Cherchneff \\& Barker 1989; Brenner \\& Barker 1989; Kurtz 1992; Cook et al.\\ 1998) and by theoretical investigations of the heating and cooling processes of PAHs in interstellar space (Allamandola, Tielens, \\& Barker 1989; Barker \\& Cherchneff 1989; d'Hendecourt et al.\\ 1989; Draine \\& Li 2001).  By far the most extensively modelled dust property is the interstellar extinction curve. Existing grain models for the diffuse interstellar medium (see Witt 2000 for a recent review) are mainly based on an analysis of extinction (Mathis, Rumpl, \\& Nordsieck 1977; Greenberg 1978; Hong \\& Greenberg 1980; Draine \\& Lee 1984; Duley, Jones, \\& Williams 1989; Mathis \\& Whiffen 1989; Kim, Martin, \\& Hendry 1994; Mathis 1996; Li \\& Greenberg 1997; Zubko 1999; Weingartner \\& Draine 2001a). The near-IR (1--5$\\mu$m), mid-IR (5--12$\\mu$m) emission spectrum along with the far-IR ($>$12$\\mu$m) continuum emission of the diffuse Galactic medium yields further insights into the composition and physical nature of interstellar dust; in particular, the PAH emission features allow us to place new constraints on the size distribution of the very small dust component.  To model this IR emission, we require realistic calculations of the excitation and deexcitation rates of small grains subject to stochastic heating by absorption of starlight, and cooling by spontaneous emission of infrared photons. We require detailed knowledge of the optical and thermal properties of interstellar dust materials to derive the probability distribution for the vibrational energy of a grain of specified size and composition.  Attempts to model the IR emission of interstellar dust have been made by various previous workers. Following the initial detection of 60 and 100$\\mu$m cirrus emission (Low et al.\\ 1984), Draine \\& Anderson (1985) calculated the IR emission from a graphite/silicate grain model with grains as small as 3\\AA\\ and argued that the 60 and 100$\\mu$m emission could be accounted for. When further processing of the IRAS data revealed stronger-than-expected 12 and 25$\\mu$m emission from interstellar clouds (Boulanger, Baud, \\& van Albada 1985), Weiland et al.\\ (1986) showed that this emission could be explained if very large numbers of 3--10\\AA\\ grains were present. A step forward was taken by D\\'{e}sert, Boulanger, \\& Puget (1990), Siebenmorgen \\& Kr\\\"{u}gel (1992), Schutte, Tielens, \\& Allamandola (1993), and Dwek et al.\\ (1997) by including PAHs as an essential grain component. Early studies were limited to the IRAS observation in four broad photometric bands, but Dwek et al.\\ (1997) were able to use DIRBE and FIRAS data. No calculations have been carried out to ascertain whether {\\it both} the PAH features {\\it and} the far-IR continuum can be quantitatively reproduced.  There has been considerable progress in both experimental measurements and quantum chemical calculations of the optical properties of PAHs (Allamandola, Hudgins, \\& Sandford 1999a; Langhoff 1996; and references therein). There is also an improved understanding of the heat capacities of dust candidate materials (Draine \\& Li 2001a, hereafter DL01) and the stochastic heating of very small grains (Barker \\& Cherchneff 1989; d'Hendecourt et al.\\ 1989; DL01) the interstellar dust size distributions (Weingartner \\& Draine 2001a, hereafter WD01a), and the grain charging processes (Weingartner \\& Draine 2001b). In the present paper we make use of these advances to model the full emission spectrum, from near-IR to submillimeter, of dust in the diffuse interstellar medium (ISM). The model is compared to IRTS, DIRBE, and FIRAS observations. We consider dust in three distinct regions: high-latitude regions ($|b|\\simgt 25^{\\rm o}$; hereafter HGL); and two regions in the Galactic plane -- the ``MIRS'' region ($44^{\\rm o}\\le l \\le 44^{\\rm o}40^{'}$, $-0^{\\rm o}40^{'}\\le b \\le 0^{\\rm o}$) for which the 4.5--11.7$\\mu$m spectrum was obtained by the mid-IR spectrometer on board IRTS; and the ``NIRS'' region ($47^{\\rm o}30^{'}\\le l \\le 48^{\\rm o}$, $|b|\\le 15^{'}$) for which the 2.8--3.9$\\mu$m spectrum was obtained by the near-IR spectrometer on board IRTS.  In the present work we consider a dust model which is a natural extension of the original graphite-silicate model for interstellar dust (Mathis, Rumpl, \\& Nordsieck 1977; Draine \\& Lee 1984, hereafter DL84). Our dust model consists of two grain types -- amorphous silicate grains and carbonaceous grains -- both with size distributions ranging from ultrasmall grains/large molecules a few angstroms in size (in the molecular domain) to large grains $\\gtsim 1\\micron$ in diameter\\footnote{% To avoid confusion, we emphasize here that ``very small grains'', ``large molecules'', and ``ultrasmall grains'' are synonymous.}. We assume that the carbonaceous grains have graphitic properties at large sizes and PAH-like properties at very small sizes. For PAH-like grains, we distinguish between the optical properties of neutral PAHs and PAH ions. The current work assumes spherical grains and therefore does not yet address issues related to polarization.  This paper is organized as follows. We first explore in \\S\\ref{sec:optc} the optical properties of interstellar dust materials, with emphasis on the absorption cross sections of both ionized and neutral PAHs from far-UV to far-IR. We then briefly summarize the heat capacities of PAHs, graphite and silicates in \\S\\ref{sec:enthalpy}. In \\S\\ref{sec:isrf} we describe the radiation field adopted for the diffuse ISM. We calculate in \\S\\ref{sec:T_eq} the equilibrium temperatures for large grains ($\\simgt 200$\\AA). In \\S\\ref{sec:T_spike} we describe the methods developed in DL01 to calculate the energy distribution functions for very small ($a\\simlt 200$\\AA) grains heated by starlight. Selected results are presented for illustration, and we demonstrate the importance of stochastic heating for grains as large as $a=200\\Angstrom$.  The adopted dust size distribution functions are summarized in \\S\\ref{sec:sizedb}. The PAH charging computation is described in \\S\\ref{sec:charging}, and \\S\\ref{sec:ir_em} details the IR emission modelling.  In \\S\\ref{sec:results} we  discuss our choices of model parameters, and compare our model with observations of infrared emission from high galactic latitudes as well as two galactic plane regions. We find our model to be in very good agreement with observations from the near-IR (including the PAH emission features) to the submillimeter. We compute emission spectra for our ``standard'' dust heated by starlight with intensities ranging from 0.3 to $10^4$ times the local value, and we tabulate the predicted intensities averaged over the SIRTF IRAC and MIPS photometric bands. The total infrared emission from our grain model is seen to be in excellent agreement with DIRBE observations, just as the albedo for our grain model is in accord with observations of the diffuse galactic light.  In \\S\\ref{sec:discussion} we show the extinction curve for our model, from the far-UV to submillimeter wavelengths, as well as tabulated values of the dust opacity. We discuss the elemental depletion, the aromatic absorption features, and the variations of the relative abundance of PAH in various regions. We also summarize the upper limits for the abundance of ultrasmall silicate grains (both amorphous and crystalline) estimated in Li \\& Draine (2001). Finally, we summarize our conclusions in \\S\\ref{sec:summary}. ",
        "conclusions": "}  Figure \\ref{fig:extcurve} compares the extinction curve obtained by our dust model with $\\bc=60$ppm ($\\bcone=45$ppm, $\\bctwo=15$ppm) with the average interstellar extinction curve and also shows the contributions of the individual dust components. It can be seen that PAHs, the dominant contributor to the 2175\\AA\\ hump, have a relatively flat behavior for $5 < \\lambda^{-1} < 8\\micron^{-1}$ and thus do not dominate the far-UV extinction. This is consistent with Greenberg \\& Chlewicki (1983) in which the 2175\\AA\\ hump was shown to correlate poorly with the far-UV extinction.  Our dust model implies a total dust mass per H nucleon $m_{\\rm dust}\\approx 1.89\\times 10^{-26}\\ {\\rm g\\ H^{-1}}$, or a gas-to-dust ratio $1.4m_{\\rm H}/m_{\\rm dust}=124.$ Let $\\Sigma_{\\rm ext}(\\lambda)\\equiv N_{\\rm H}^{-1}\\int (dn/da) C_{\\rm ext}(a,\\lambda)da$ be the total extinction cross section per H nucleon. The mass absorption coefficient for the dust is just $\\kappa_{\\rm abs}(\\lambda) = (1-\\albedo)\\Sigma_{\\rm ext}(\\lambda)/m_{\\rm dust} = 0.4\\ {\\rm ln}10\\ (1-\\albedo) [A(\\lambda)/N_{\\rm H}]/m_{\\rm dust} = 4.88\\times 10^{25} (1-\\albedo) [A(\\lambda)/N_{\\rm H}]\\ {\\rm cm^2\\ g^{-1}}$. In Table \\ref{tab:kext} we tabulate $\\kappa_{\\rm abs}$ at selected wavelengths. Note that these numbers are for dust in the average diffuse ISM, with $R_V \\equiv A_V/E(B-V)\\approx 3.1$. In diffuse regions with $R_V > 3.1$ the extinction per H is reduced at $\\lambda < 0.9\\micron$ due to changes in the grain size distribution (see WD01), but the extinction does not appear to be affected at $\\lambda \\gtsim 0.9\\micron$ (Cardelli, Clayton, \\& Mathis 1989). In dark clouds ice mantle formation adds additional absorption features, most prominently the 3.1$\\micron$ O-H stretch in H$_2$O, but for $\\lambda \\gtsim 10\\micron$ the opacities in Table \\ref{tab:kext} probably remain a good approximation even for molecular regions with densities up to $\\sim 10^6 \\cm^{-3}$.  \\begin{deluxetable}{rccccc} \\tabletypesize{\\footnotesize} \\tablewidth{0pt} \\tablecaption{Extinction and absorption properties for dust in the diffuse ISM.\\tablenotemark{a}\\newline NOTE: values of $\\Sigma_{\\rm ext}$ and $\\kappa_{\\rm abs}$ in ApJ Table 6 are too small by a factor of 1.086 . \\label{tab:kext}} \\tablehead{ \\colhead{band}& \\colhead{$\\lambda$}& \\colhead{albedo}& \\colhead{$g\\equiv\\langle\\cos\\theta\\rangle$ \\tablenotemark{b}}& \\colhead{$\\Sigma_{\\rm ext}$ \\tablenotemark{c}}& \\colhead{$\\kappa_{\\rm abs}$ \\tablenotemark{d}} \\\\ \\colhead{}& \\colhead{($\\mu$m)}& \\colhead{}& \\colhead{}& \\colhead{($\\cm^{2}\\ {\\rm H}^{-1}$)}& \\colhead{($\\cm^{2}\\g^{-1}$)} } \\startdata & 0.0100\t& 0.47\t& 1.00 &$7.11\\times 10^{-22}$ & $2.00\\times 10^{4}$ \\\\ & 0.0488\t& 0.30\t& 0.86 &$1.52\\times 10^{-21}$ & $5.66\\times 10^{4}$ \\\\ Ly edge  &0.0912& 0.24\t& 0.73 &$2.51\\times 10^{-21}$ & $1.01\\times 10^{5}$ \\\\ & 0.100\t& 0.27\t& 0.72 &$2.28\\times 10^{-21}$ & $8.83\\times 10^{4}$ \\\\ Ly$\\alpha$&0.1215&0.32\t& 0.73 &$1.73\\times 10^{-21}$  & $6.20\\times 10^{4}$ \\\\ & 0.150\t& 0.38\t& 0.70 &$1.38\\times 10^{-21}$  & $4.56\\times 10^{4}$ \\\\ & 0.220\t& 0.42\t& 0.56 &$1.65\\times 10^{-21}$  & $5.05\\times 10^{4}$ \\\\ & 0.300\t& 0.58\t& 0.57 &$1.00\\times 10^{-21}$ & $2.25\\times 10^{4}$ \\\\  U & 0.365\t& 0.62\t& 0.58 &$8.28\\times 10^{-22}$ & $1.68\\times 10^{4}$ \\\\ B & 0.440\t& 0.65\t& 0.57 &$6.73\\times 10^{-22}$ & $1.26\\times 10^{4}$ \\\\ V & 0.550\t& 0.67\t& 0.54 & $5.06\\times 10^{-22}$ & 8960 \\\\ R & 0.700\t& 0.66\t& 0.48 &$3.59\\times 10^{-22}$  & 6460 \\\\ I & 0.900\t& 0.63\t& 0.40 &$2.47\\times 10^{-22}$  & 4820 \\\\ J & 1.22\t& 0.58  & 0.29 &$1.53\\times 10^{-22}$  & 3440\\\\ H & 1.63\t& 0.51  & 0.21 &$9.76\\times 10^{-23}$  & 2510\\\\ K & 2.20\t& 0.43  & 0.13 &$5.92\\times 10^{-23}$  & 1780\\\\ L & 3.45\t& 0.28  & 0.0051  &$2.57\\times 10^{-23}$ & 984\\\\ & 3.60\t& 0.26  & -0.0037 &$2.37\\times 10^{-23}$ & 926\\\\ & 4.50\t& 0.18  & -0.037 &$1.54\\times 10^{-23}$\t& 670\\\\ M & 4.80 \t& 0.16  & -0.043 &$1.37\\times 10^{-23}$\t& 611\\\\ & 5.80\t& 0.099 & -0.052 &$1.02\\times 10^{-23}$ & 488\\\\ & 8.00\t& 0.016 & -0.050 &$1.54\\times 10^{-23}$\t& 803\\\\ & 9.70\t& 0.003 & -0.036 &$4.39\\times 10^{-23}$\t& 2320\\\\ N & 10.6\t& 0.003 & -0.032 &$3.36\\times 10^{-23}$\t& 1770\\\\ & 12.0\t& 0.003 & -0.029 &$1.93\\times 10^{-23}$\t& 1020\\\\ Q & 21.0\t& $5\\times 10^{-4}$ & -0.024 &$1.30\\times 10^{-23}$ & 693\\\\ & 25.0\t& $3\\times 10^{-4}$ & -0.024 &$9.29\\times 10^{-24}$ & 493\\\\ & 60.0\t& $1\\times 10^{-4}$ &$< 0.01$ &$1.91\\times 10^{-24}$ & 101\\\\ & 70.0\t& $1\\times 10^{-4}$ &$< 0.01$ &$1.38\\times 10^{-24}$ & 73.1\\\\ & 100\t\t& $< 10^{-4}$ & $< 0.01$ &$6.40\\times 10^{-25}$ & 34.0 \\\\ & 140\t\t& $< 10^{-4}$ & $< 0.01$ &$3.31\\times 10^{-25}$ & 17.6 \\\\ & 160\t\t& $< 10^{-4}$ & $< 0.01$ &$2.44\\times 10^{-25}$ & 13.0 \\\\ & 240\t\t& $< 10^{-4}$ & $< 0.01$ &$1.06\\times 10^{-25}$ & 5.61 \\\\ & 450\t\t& $< 10^{-4}$ & $< 0.01$ &$2.80\\times 10^{-26}$ & 1.49 \\\\ & 850\t\t& $< 10^{-4}$ & $< 0.01$ &$8.83\\times 10^{-27}$ & 0.468 \\\\ & 1350\t& $< 10^{-4}$ & $< 0.01$ &$4.05\\times 10^{-27}$ & 0.215 \\\\ & 2000\t& $< 10^{-4}$ & $< 0.01$ &$2.10\\times 10^{-27}$ & 0.111 \\\\ & 3000\t& $< 10^{-4}$ & $< 0.01$ &$1.05\\times 10^{-27}$ & 0.0559 \\\\ \\hline \\enddata \\vspace*{-0.3cm} \\tablenotetext{a}{An expanded version of this table is available at {\\tt http://www.astro.princeton.edu/$\\sim$draine/dust/dustmix.html}} \\tablenotetext{b}{Average of $\\cos\\theta$ for scattered light, where $\\theta$ is the scattering angle.} \\tablenotetext{c}{$\\Sigma_{\\rm ext}$ is the dust extinction cross section per H nucleon.} \\tablenotetext{d}{$\\kappa_{\\rm abs}$ is the absorption opacity of the dust alone; the opacity of the gas+dust mixture is $\\kappa_{\\rm abs}/125$. To within $\\pm$10\\%, $\\kappa_{\\rm abs}\\approx 3.17\\times 10^5 \\left(\\lambda/\\mu {\\rm m}\\right)^{-2} ({\\rm cm^2\\ g^{-1}})$ for $20 \\!<\\! \\lambda \\!<\\! 700\\micron$, and $\\kappa_{\\rm abs} \\approx 3.89\\times 10^4\\ \\left(\\lambda/\\mu {\\rm m}\\right)^{-1.68} {\\rm cm^2\\ g^{-1}}$ for $700 \\!<\\! \\lambda \\!<\\! 10^4\\micron$.} \\end{deluxetable}   Our dust model requires a total of 254ppm of C and 48 ppm of Si in grains, both of which are significantly higher than the maximum available for the dust according to the recent work of Snow \\& Witt (1995), who suggest that the elemental abundances in the ISM are just 60\\%--70\\% of the solar value\\footnote{Although the present model uses more silicon than other existing models, the predicted silicate optical depth $\\Delta \\tau_{\\rm 9.7\\mu m}/N_{\\rm H} \\approx 3.3\\times 10^{-23}\\ {\\rm cm^2/H}$ is not inconsistent with observations ($\\Delta \\tau_{\\rm 9.7\\mu m}/N_{\\rm H} \\approx 2.8-5.9\\times 10^{-23}\\ {\\rm cm^2/H}$; see Draine 1989b, Mathis 1998).}. However, Weingartner \\& Draine (2001a) have argued that it is difficult to see how the observed interstellar extinction can be explained if interstellar abundances are appreciably subsolar; they argue that the case for subsolar abundances is not compelling. Dwek (1997) has argued that the composite dust model (Mathis 1996) -- which aims at making economical use of dust materials -- is still short of carbon when PAHs are included to account for the mid-IR emission.  Our model has an absorption feature at 6.2$\\mu$m due to the PAH C-C stretching mode (see Figure \\ref{fig:extcurve}). The integrated optical depth of the 6.2$\\mu$m feature relative to the silicate optical depth is $(\\int \\Delta\\tau d\\lambda^{-1})_{6.2} /\\Delta\\tau_{9.7} \\approx 2.9\\cm^{-1}$. The aromatic 6.2$\\mu$m absorption feature has been detected in several objects including both local sources and Galactic center sources (Schutte et al.\\ 1998; Chiar et al.\\ 2000). On average, the observed integrated strength of the 6.2$\\mu$m band per unit silicate band depth is about 2.2cm$^{-1}$ with no significant difference between the Galactic sources and the local sources (Schutte et al.\\ 1998). Our model is in good agreement with this determination.  In addition to the 6.2$\\mu$m feature, our model also predicts a weak narrow absorption feature at 3.3$\\mu$m with $\\int \\Delta\\tau_{3.3} d\\lambda^{-1}/\\Delta\\tau_{9.7} \\approx 0.50\\cm^{-1}$. This feature has been detected in the Galactic center source GCS 3 with $(\\int \\Delta\\tau d\\lambda^{-1})_{3.3}/\\Delta\\tau_{9.7} \\approx 0.34\\cm^{-1}$ (Chiar et al.\\ 2000). It has also been seen in absorption in some heavily extincted molecular cloud sight lines (Sellgren et al.\\ 1995; Brooke, Sellgren, \\& Geballe 1999). The 7.7, 8.6, and 11.3$\\mu$m features are hidden by the much stronger 9.7$\\mu$m silicate feature, and therefore will be difficult to observe as absorption features.  The aromatic absorption bands allow us to place constraints on the PAH abundance if the PAH band strengths are known. On the other hand, one can also infer information on the PAH band strengths from the aromatic absorption bands given knowledge of the PAH abundance derived from other considerations such as the 2175\\AA\\ hump and/or the 3--12$\\mu$m PAH emission bands. For example, for $\\bc$=60ppm in PAHs (the maximum allowed by the 2175\\AA\\ hump), and assuming $a_{\\xi}\\approx50\\Angstrom$, enhancement of the 6.2$\\micron$ band by a factor $E_{6.2} > 3$ would result in a 6.2$\\mu$m absorption feature which would be stronger than observed; conversely, if we had not enhanced the 6.2$\\mu$m feature strength by a factor $E_{6.2}$=3, then the 6.2$\\mu$m absorption feature strength in the model would be weaker than observed. However, this depends on the poorly-constrained value of $a_{\\xi}$, which determines the radius where carbonaceous grains make the transition from being PAH-like to becoming graphite-like. It is clear that $(\\int \\Delta\\tau d\\lambda^{-1})_{6.2}$ would increase (decrease) as $a_{\\xi}$ is increased (decreased) (e.g., the $a_{\\xi}=100$\\AA\\ model results in $(\\int \\Delta\\tau d\\lambda^{-1})_{6.2} /\\Delta\\tau_{9.7} \\approx 4.2\\cm^{-1}$). Therefore, until we have better understanding of the PAH-graphite transition, we should not over-state the constraints of the aromatic absorption bands on the PAH band strengths. However, the fact that our model is consistent with the 3.3, 6.2$\\mu$m absorption bands, the 2175\\AA\\ hump, and the 3--12$\\mu$m emission indicates that the choice of $a_{\\xi}$=50\\AA\\ is reasonable.  We must also stress that the strengths of these PAH absorption bands could vary with physical conditions due to changes in the PAH ion fraction $\\fion(a)$. In regions with increased $\\fion$, the 6.2$\\mu$m absorption feature would be strengthened and the 3.3$\\mu$m feature would be weakened (see Table \\ref{tab:drude_parameters}). The fact that there is no significant variation between the Galactic center sources and the local sources for the 6.2$\\mu$m absorption strength relative to the silicate optical depth (Schutte et al.\\ 1998) suggests that these regions probably have similar radiation field-to-electron density ratios. This is also indicated by the $P(7.7\\mu {\\rm m})/P(11.3\\mu {\\rm m})$ data of Chan et al.\\ (2001) (see their Figure 2c).  Significant variations in the relative abundances of PAHs to large grains are expected in different environments due to depletion (e.g., accretion of PAHs into grain mantles, agglomeration of PAHs and other very small grains to form composite grains), production (e.g. fragmentation of graphitic grains, hydrogenated amorphous carbon [Scott, Duley, \\& Pinho 1997] or organic refractory materials [Greenberg et al.\\ 2000]), or drift of large grains relative to small grains due to anisotropic starlight (Weingartner \\& Draine 2001c). Such local variations in PAH abundance are indicated by the large cloud-to-cloud variations in mid-IR emission (Verter et al.\\ 2000) as well as the lower UV extinction in dense regions.  Ultrasmall silicate grains are excluded as a major component since the IRTS and ISO spectroscopic observations of the diffuse ISM show no evidence of the Si-O 9.7$\\mu$m emission feature which would be expected for transiently-heated tiny silicate dust (Mattila et al.\\ 1996; Onaka et al.\\ 1996). However, the presence of these particles cannot be ruled out since a weak 9.7$\\mu$m emission feature may be hidden by the strong PAH features at 8.6$\\micron$ and 11.3$\\micron$. On the basis of the observed IR emission spectrum for the diffuse ISM, the observed ultraviolet extinction curve, and the 10$\\mu$m silicate absorption profile, Li \\& Draine (2001) have obtained upper limits on the abundances of ultrasmall ($a\\ltsim15\\Angstrom$) amorphous and crystalline silicate grains. It is found that, contrary to previous work ($\\simlt 1\\%$ of solar silicon abundance [D\\'esert et al.\\ 1986]), as much as $\\sim$10\\% of interstellar Si could be in $a\\ltsim 15\\Angstrom$ silicate grains without violating observational constraints. Not more than $\\sim$5\\% of the Si can be in crystalline silicates (of any size).  }  We have modeled the infrared emission spectra from dust in the diffuse ISM of the Milky Way, including the prominent emission features at 3.3, 6.2, 7.7, 8.6 and 11.3$\\mu$m and the continuum emission from near-IR to submillimeter. We have compared our model to the observed emission from high galactic latitude regions as well as from regions in the Galactic plane.  The dust model adopted here consists of a mixture of amorphous silicate grains and carbonaceous grains, both having a wide size distribution ranging from large grains $\\sim$1$\\micron$ in diameter. down to molecules containing tens of atoms. We assume that the carbonaceous grains have graphitic properties at radii $a \\gtrsim 50$\\AA, and polycyclic aromatic hydrocarbon (PAH)-like properties at the very small sizes (see \\S\\ref{sec:optc_gra}) which account for the 3.3, 6.2, 7.7, 8.6 and 11.3$\\mu$m emission features seen in a wide range of objects.  Based on recent progress made in both laboratory and theoretical studies of PAH absorption spectra, and guided by the astronomical observations, we have constructed ``astronomical'' absorption cross sections for both neutral and ionized PAHs from the far ultraviolet to the infrared (see \\S\\ref{sec:optc_pah}, Table \\ref{tab:drude_parameters}, and Appendix \\ref{sec:appendix}) including dependence on grain size (which affects the onset of the absorption edge at long wavelength [see \\S\\ref{app:visible_to_uv}]), H/C ratio, and charge state. Detailed modelling indicates that the feature strengths of the 7.7, 8.6$\\mu$m bands need to be increased by a factor $\\sim$2, and a factor $\\sim$3 for the 6.2$\\mu$m band (see \\S\\ref{sec:optc_pah} and compare Figures \\ref{fig:irts_chi2}a and \\ref{fig:irts_e1_h2c}a). In addition, a small amount of continuum absorption needs to be added to the PAH absorption spectrum to account for the near-IR continuum emission at 1--5$\\mu$m underlying the 3.3$\\mu$m C-H stretching feature detected in various sources (see \\S\\ref{sec:pahircont}). We represent this continuum opacity by adding about 1\\% of the opacity of graphite. We have also slightly modified the far-infrared and submillimeter emissivity of ``astronomical silicate'' in order to explicitly demonstrate that the observed emission from the interstellar medium can be reproduced using reasonable properties for carbonaceous and silicate grains.  We model the heat capacities of carbonaceous and silicate grains using vibrational models developed by DL01. The thermal-discrete method developed in DL01 is used to calculate the energy (temperature) distribution functions for transiently heated very small grains (see \\S\\ref{sec:T_spike}). We find that for grains heated by starlight with $\\chiMMP=1$, the energy distribution function is quite broad even for grains as large as $a=200\\Angstrom$, resulting in stronger emission at short wavelengths than would be computed using the ``equilibrium temperature'' $\\bar{T}$. The PAH charging computation is performed using the technique developed by Weingartner \\& Draine (2001b) (see \\S\\ref{sec:charging}).  We adopt the size distributions for large grains (see \\S\\ref{sec:sizedb}) derived from modeling the interstellar extinction curve (Weingartner \\& Draine 2001a). We find that if the very small carbonaceous grain population includes two log-normal size distributions characterized by $a_{01}=3.5$\\AA, $a_{02}=30$\\AA, and widths $\\sigma=0.4$ for the PAH size distributions (with about 60ppm of C [relative to H] contained in PAHs; 45ppm in the former, and 15ppm in the latter), then the overall grain model is able to closely reproduce the observed emission of the diffuse ISM from near-IR to submillimeter (see \\S\\ref{sec:results}) while at the same time reproducing the observed extinction curve. Our dust model predicts a pronounced absorption feature at 6.2$\\mu$m and a much weaker feature at 3.3$\\mu$m due to PAHs with optical depths consistent with observations (see \\S\\ref{sec:discussion}).  We will show elsewhere (Draine \\& Li 2000b) that the very same ultrasmall grain component (i.e. PAHs) invoked here to account for the 3--60$\\mu$m {\\it vibrational} emission will produce {\\it ``rotational''} emission which can quite naturally explain the observed 10--100GHz ``anomalous'' microwave emission from the interstellar medium, confirming the original estimates by Draine \\& Lazarian (1998a,b).  While real interstellar grains are no doubt more complex than the simple model presented here, this extension of the original graphite-silicate grain model (Mathis, Rumpl, \\& Nordsieck 1977; Draine \\& Lee 1984) to explicitly include a PAH population leads to a well-defined grain model which is in excellent agreement with available observational constraints, ranging from the observed extinction curve to the observed infrared emission."
    },
    "0011/astro-ph0011069_arXiv.txt": {
        "abstract": "We determine the angular power spectrum $C_l$ of the Edinburgh/Durham Southern Galaxy Catalog (EDSGC) and use this statistic to constrain cosmological parameters.  Our methods for determining $C_l$ and the parameters that affect it are based on those developed for the analysis of cosmic microwave background maps.  We expect them to be useful for future surveys. Assuming flat cold dark matter models with a cosmological constant (constrained by {\\it COBE}/DMR and local cluster abundances) and a scale--independent bias $b$, we find acceptable fits to the EDSGC angular power spectrum with $1.11 < b < 2.35$ and $0.2 < \\Omega_m < 0.55$ at 95\\% confidence.  These results are not significantly affected by the ``integral constraint'' or extinction by interstellar dust, but may be by our assumption of Gaussianity. ",
        "introduction": "Over the next decade, the quantity and quality of galaxy survey data will improve greatly because of a variety of new survey projects underway including the Sloan Digital Sky Survey (SDSS; see York et al.\\ 2000).  However, most of the galaxies in such surveys will {\\it not} have spectroscopically--determined redshifts; therefore, the study of their angular correlations will be highly profitable for our understanding of the large--scale structure of the Universe.  The primary purpose of this paper is to consider an analysis approach which is likely to be useful for deriving cosmological constraints from these larger surveys.  In particular, we use methods that have become standard in the analysis of cosmic microwave background (CMB) anisotropy maps, such as those from {\\it Boomerang} (de Bernardis et al.\\ 2000; Lange et al.\\ 2000) and {\\it Maxima}-I (Hanany et al.\\ 2000, Balbi et al.\\ 2000).  Estimation of the two--point angular correlation function $w(\\theta)$ from galaxy surveys without redshfit information has a long history.  Early work (Peebles \\& Hauser 1974; Groth \\& Peebles 1977) found the angular correlation function to vary as $w(\\theta)=\\theta^{1-\\gamma}$ with $\\gamma= 1.77$ and a break at scales larger than $\\sim 9\\hmpc$. The advent of automated surveys, such as the APM galaxy survey (Maddox et.\\ al.\\ 1990) and Edinburgh/Durham Southern Galaxy Catalogue (EDSGC; Collins, Nichol \\& Lumsden 1992) enabled a much more accurate determination of $w(\\theta)$, as each survey contained angular positions for over a million galaxies.  One way to compare the measured angular correlation function with theoretical predictions is to invert $w(\\theta)$ to obtain the three--dimensional power spectrum $P(k)$. This requires inverting Limber's equation (Limber 1953).  Baugh \\& Efstathiou (1993, 1994) and Gazta\\~{n}aga \\& Baugh (1998) used Lucy's algorithm (Lucy 1974) to do the inversion, while Dodelson \\& Gazta\\~{n}aga (2000) used a Bayesian prior constraining the smoothness of the power spectrum.  Eisenstein \\& Zaldarriaga (2000) used a technique based on singular-value decomposition to get $P(k)$ from $w(\\theta)$.  They point out that, once the correlations in the inverted power spectra are included, the uncertainties on cosmological parameters from the APM are significantly weakened.  Our analysis is a three-step process, similar to what is done with CMB data sets (Tegmark 1997; Bond, Jaffe \\& Knox 1998; Bond, Jaffe \\& Knox 2000). The first step is the construction of a pixelized map of galaxy counts, together with its noise properties.  The second step is the determination of the angular power spectrum $C_l$ of the map using likelihood analysis, together with window functions and a covariance matrix.  In the final step, we compare our observationally--determined $C_l$ to the $C_l$ predicted for a given set of parameters in order to get constraints on those parameters.  We assume that the errors in $C_l$ are log--normally distributed.  The angular power spectrum $C_l$ is a useful intermediate step on this road from galaxy catalog to parameter constraints. Estimates of the angular power spectrum, together with a description of the uncertainties, can be viewed as a form of data compression.  One has converted the $\\sim\\!1$ million EDSGC galaxies (for example) into a handful of power spectrum constraints, together with window functions and covariance matrices.  Thus if one wishes to make other assumptions about bias and cosmological parameters than we have done here and determine the resulting constraints, one can do so {\\it without} having to return to the cumbersome galaxy catalog.  We use $C_l$ instead of its historically--preferred Legendre transform $w(\\theta)$ for several reasons. First, the error matrix structure is much simpler: $\\langle \\delta C_l \\delta C_{l'} \\rangle$ is band-diagonal and becomes diagonal in the limit of full--sky coverage, whereas $\\langle \\delta w(\\theta) \\delta w(\\theta') \\rangle$ is much more complicated and does not become diagonal even in the full--sky limit. Second, the relation between $C_l$ and the corresponding three--dimensional statistic $P(k)$ is simpler than that between $w(\\theta)$ and $P(k)$ (or its Fourier transform $\\xi(r)$ (Baugh \\& Efstathiou 1994)).  We use likelihood analysis to determine $C_l$ because the likelihood is a fundamental statistical quantity.  The likelihood is the probability of the data given $C_l$, which by Bayes' theorem is proportional to the probability of $C_l$ given the data.  Another advantage of likelihood analysis is that, as explained below, it allows for straightforward control of systematic errors (due to, e.g., masking) via modifications of the noise matrix.  Only on sufficiently large scales do we expect the likelihood function to be a Gaussian that depends only on $C_l$ and not any higher-order correlations.  We therefore restrict our analysis to $l$ values less than some critical value.  On small scales the likelihood function becomes much more complicated and its form harder to predict {\\it a priori}.  Mode-mode coupling due to nonlinear evolution leads to departures of the $C_l$ covariance matrix from band-diagonal. Therefore some of the advantages of likelihood analysis and the angular power spectrum are lost on smaller scales where other techniques may be superior.  The Gaussianity assumption is perhaps the weakest point of the approach outlined here.  Below, we briefly discuss how the analysis can be improved in this regard with future data sets.  The EDSGC, with over a million galaxies and covering over 1000 sq.\\ degrees, offers us an excellent test bed for applying our algorithms (Nichol, Collins \\& Lumsden 2000).  We convert this catalog into a pixelized map and determine its angular power spectrum together with window functions and covariance matrix. As an illustrative application of the angular power spectrum, we constrain a scale--independent bias parameter $b$ and the cosmological constant density parameter $\\Omega_\\Lambda$ in a {\\it COBE}--normalized $\\Lambda$CDM model with zero mean spatial curvature.  Our constraints on the bias are improved by including constraints on the amplitude of the power spectrum derived from number densities of low--redshift massive clusters of galaxies (Viana \\& Liddle 1999, hereafter VL99; also see Pierpaoli, Scott \\& White 2000).  These number densities are sensitive to the amplitude of the matter power spectrum calculated in linear perturbation theory, near the range of length scales probed by the EDSGC.  The angular power spectrum of the APM catalog was previously estimated by Baugh \\& Efstathiou (1994) though not via likelihood analysis. Very recently, Efstathiou \\& Moody (2000) have applied the same techniques we do here to estimating $C_l$ for the APM survey.  Their approach differs from ours in how they constrain cosmological parameters. Instead of projecting the theoretical three--dimensional power spectra $P(k)$ into angular power spectra, they transform their $C_l$ constraints into (highly correlated) constraints on $P(k)$ and then compare to theoretical $P(k)$.  We expect the analysis methods presented here to be useful for other current and future data sets---even those with large numbers of measured redshifts.  For example, the Sloan Digital Sky Survey (SDSS) will spectroscopically determine the redshifts of a million galaxies, but there will be about one hundred times as many galaxies in the photometric data, without spectroscopic redshifts.  One can generalize the methods presented here to analyze sets of maps produced from galaxies in different photometric redshift slices.  In \\S \\ref{sec:likelihood}, we review likelihood analysis and the use of the quadratic estimator to iteratively find the maximum of the likelihood function.  In \\S \\ref{sec:Cl} we describe our calculation of $P(k)$ and its projection to $C_l$.  In \\S \\ref{sec:extract} we show how to compare the calculated $C_l$ to the measured $C_l$ in order to determine parameters.  In \\S \\ref{sec:edsgc} we apply our methods to the EDSGC, and discuss some possible sources of systematic error in \\S \\ref{sec:sys}. This is followed by a discussion of our results in \\S \\ref{sec:discussion} and a brief conclusion in \\S \\ref{sec:conclusions}.  An appendix outlines the derivation of the projection of $P(k)$ to $C_l$. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a general formalism to analyze galaxy surveys without redshift information.  We pixelize the galaxy counts on the sky, and then, using the quadratic estimator algorithm, extract the angular power spectrum -- a procedure already in use in CMB data analysis. Just like in the CMB case, one effectively converts complex information contained in the experiment (in this case, locations of several hundred thousand galaxies) into a handful of numbers -- the angular power spectrum. One can then use the angular power spectrum for all subsequent analyses.  We apply this method to the EDSGC survey. We compute the angular power spectrum of EDSGC, and combine it with {\\it COBE}/DMR and cluster constraints to obtain constraints on cosmological parameters. Assuming flat $\\Lambda$CDM models with constant bias between galaxies and dark matter, we get $1.11 < b < 2.35$ and $0.2 < \\Omega_m < 0.55$ at 95\\% confidence.  One advantage of our formalism is that it does {\\it not} require galaxy redshifts, but only their positions in the sky. This should make it useful for surveys with very large number of galaxies, only a fraction of which will have redshift information. For example, the ongoing SDSS is expected to collect about one million galaxies with redshift information, but also a staggering one hundred million galaxies with photometric information only. Using the techniques presented in this paper, one will be able to convert that information into the angular power spectrum, which can then be used for various further analyses."
    },
    "0011/astro-ph0011543_arXiv.txt": {
        "abstract": "We study a broad-band X-ray observation by \\sax\\/ of the broad-line radio galaxy 3C 120. The primary X-ray continuum is well described by a power law with the photon index of $\\Gamma\\sim 1.85\\pm 0.05$ and a high-energy break or a cutoff. The obtained e-folding energy of $\\sim 100$--300 keV corresponds to $kT\\sim 50$ keV in thermal-Compton models. A soft X-ray excess is found at a statistical significance of $\\ga 98\\%$. Our physical best-fit model of the excess is optically-thin emission from an extended plasma (with a luminosity of $\\sim 2\\%$ of the total X-ray luminosity), which interpretation is supported by an independent finding of an extended soft X-ray halo in 3C 120. We find a moderate Compton reflection component together with a moderately broad Fe K$\\alpha$ line with an equivalent width ($\\sim 10^2$ eV) fully consistent with the strength of reflection, indicating their origin in an optically-thick accretion disk extending relatively close to the central black hole. We observe strong spectral variability during our 2-day observation with $\\Gamma$ varying from $\\sim 1.7$ to $\\sim 2$ and correlated with the soft X-ray flux. The variability is very well modeled by thermal Comptonization in a hot plasma irradiated by a variable UV flux. Both a hot patchy corona outflowing with a mildly relativistic velocity away from a cold disk and a hot inner accretion flow overlapping with the cold disk are viable geometries. The presence of an outflow in the first case and an overlap between the hot and cold phases in the second case are required by energy balance and the observed strength of Compton reflection. ",
        "introduction": "\\label{s:intro}  3C 120 is a broad-line radio galaxy (BLRG) at $z=0.033$ exhibiting characteristics intermediate between those of FR I galaxies and BL Lacs (Urry \\& Padovani 1995). It is also classified as a Seyfert 1 (Burbidge 1967), although its optical morphology is complex, see Grandi et al.\\ (1997).  It is variable in all energy bands, and shows superluminal motion in the VLBI core (e.g., Walker 1997).  In 1994, 3C 120 has been observed by \\asca, and since then the strength and the width of its Fe K$\\alpha$ has been a subject of a considerable controversy. Reynolds (1997) found the line to be extremely strong, with an equivalent width of $\\wfe\\approx 1$ keV, as well as extremely broad, with a Gaussian width of $\\sfe\\approx 2$ keV. The physical origin of such a line is unclear. A slightly weaker line, with $\\wfe\\approx 0.5$ keV and $\\sfe\\sim 1$ keV, was found in the same data by Grandi et al.\\ (1997). Both results were obtained assuming the underlying continuum was a single power law in the entire energy range observed by \\asca, and Grandi et al.\\ (1997) mentioned the possibility that the unusual line parameters  were caused by that assumption not being valid. Indeed, Wo\\'zniak et al.\\ (1998, hereafter W98) found that if the \\asca\\/ data (with an updated calibration of 1997) were fitted in the 4--10 keV range assuming the presence of Compton reflection with a reflector solid angle of $\\Omega$ between 0 and $2\\pi$, the parameters of the line became much more physically plausible, with the best-fit values of $\\wfe\\approx 60$--110 eV, the peak energy, $\\efe \\approx 6.4$ keV, and $\\sfe\\simeq 0.32$ keV and 0.20 keV assuming $\\Omega=0$ and $2\\pi$, respectively. The line in the \\asca\\/ data was again claimed to be very strong and broad by Sambruna, Eracleous \\& Mushotzky (1999), who obtained $\\wfe=510_{- 90}^{+70}$ eV and $\\sfe=0.89^{+0.17}_{- 0.15}$ keV, virtually identical with the values of Grandi et al.\\ (1997).  This issue has recently been resolved thanks to results of \\xte\\/ and \\sax. Eracleous, Sambruna \\& Mushotzky (2000, hereafter E00) have found the line in the \\xte\\/ data ($\\geq 4$ keV) to be weak with the parameters consistent with those of W98, namely $\\wfe=90^{+30}_{- 20}$ eV and $\\sfe<1.2$ keV. Similarly, Grandi (2000, hereafter G00) have found $\\wfe=60^{+80}_{- 40}$ eV and $\\sfe=0.27_{-0.27}^{+0.40}$ keV in a preliminary analysis of the \\sax\\/ data.  Since \\ginga\\/ did not observe 3C 120, the actual strength of Compton reflection had been unknown until the results of E00 and G00, who found $\\Omega/2\\pi=0.4^{+0.4}_{- 0.1}$ and  $0.7^{+0.4}_{-0.4}$, respectively. These values of the solid angle are compatible (George \\& Fabian 1991; \\.Zycki \\& Czerny 1994) with the presence of a modest Fe K$\\alpha$ line with the parameters of W98, E00 and G00.  Here, we present a comprehensive analysis of the observation of 3C 120 by \\sax, which combines the broad-band hard X-ray coverage of \\xte\\/ with a high energy resolution and an energy range in soft X-rays comparable to that of \\asca. We address the issues of the parameters of the line and the continuum, their variability, radiative processes and geometry of the source, and the relationship of 3C 120 to other BLRGs as well as Seyferts. ",
        "conclusions": "\\label{s:con}  Our \\sax\\/ data clearly show a continuum due to reflection of an X-ray power law with $\\Gamma\\sim 1.8$--1.9 from a cold material. The significance of the presence of reflection is $> 99.99\\%$ and the solid angle of the reflector is $\\la 2\\pi$ ($\\Omega/2\\pi= 0.56_{-0.20}^{+0.44}$ for the e-folded power-law incident continuum). The reflection is accompanied by a fluorescent Fe K$\\alpha$ line with an equivalent width ($\\sim 80$ eV) fully consistent with originating entirely from the reflector. The line is moderately broad, and combining our spectral constraints with those from \\asca\\/ and \\xte\\/ (which data show an overall agreement with those of \\sax) yields $\\sfe\\simeq 0.2$--0.3 keV. If this line comes from a Keplerian disk in a Schwarzschild metric, the inner radius of the emitting region is $\\ga 10GM/c^2$.  We find the presence of a high-energy break or cutoff in the \\sax\\/ data. We can fit the data by either an e-folded power law or thermal Comptonization, obtaining $E_{\\rm c}= 150_{-30}^{+230}$ keV, $kT\\sim 40$--130 keV, respectively. In the latter model, $\\tau\\sim 0.6$--3. We also detect a relatively weak (but significant at a $\\sim 98\\%$ confidence) soft excess. It is best-fitted by emission from an extended thermal plasma at $kT_{\\rm s}\\simeq 1$ keV, which model is supported by an independent finding of an extended soft X-ray halo from 3C 120.  We find strong spectral variability during the 2-day observation, corresponding to  $\\Gamma$ varying from $\\sim 1.7$ to $\\sim 2$ and correlated with the soft X-ray flux. The variability is very well modeled by thermal Comptonization in a hot plasma with a constant $\\tau$ which temperature adjusts to a variable flux of irradiating UV photons.  The 3C 120 data obey the general correlation between the power-law slope and the strength of Compton reflection found in Seyferts and black-hole binaries. Interpretation of this correlation involves feedback between a thermal Comptonizing plasma and a cold medium, yielding support for our thermal-Compton spectral model. Our data are consistent with two geometries: an outflow above the surface of a disk, or a hot inner disk surrounded by an overlapping cold disk. Both homogeneous and static patchy corona models can be ruled out.  We confirm (at a very high confidence level) that BLRGs have harder X-ray spectra, less Compton reflection and weaker Fe K$\\alpha$ lines than those average in radio-quiet Seyferts. However, the distributions of those parameters in BLRGs and radio-quiet Seyferts are not disjoint; the former occupy a part of the parameter space occupied by the latter."
    },
    "0011/astro-ph0011296_arXiv.txt": {
        "abstract": "We report on the discovery of variable circular polarization in the SW Sex star LS Pegasi. The observed modulation has an amplitude of $\\approx 0.3~\\%$ and a period of 29.6 minutes, which we assume as the spin period of the magnetic white dwarf. We also detected periodic flaring in the blue wing of H$_\\beta$, with a period of 33.5 minutes. The difference between both frequencies is just the orbital frequency, so we relate the 33.5-min modulation to the \\emph{beat} between the orbital and spin period. We propose a new accretion scenario in SW Sex stars, based on the shock of the disk-overflown gas stream against the white dwarf's magnetosphere, which extends to the corotation radius. From this geometry, we estimate a magnetic field strength of $B_\\mathrm{1}\\sim 5$--15 MG. Our results indicate that magnetic accretion plays an important role in SW Sex stars and we suggest that these systems are probably Intermediate Polars with the highest mass accretion rates. ",
        "introduction": "SW Sex systems are a sub-class of nova-like cataclysmic variables (CVs) with peculiar spectroscopic behavior \\citep{thorstensen91}. They exhibit strong single-peaked Balmer, \\ion{He}{1} and \\ion{He}{2} emission lines which, with the exception of \\ion{He}{2}, remain largely unobscured during primary eclipse (suggesting emission in vertical structures). The radial velocity curves of Balmer and \\ion{He}{1} lines show significant phase lags relative to photometric minimum, and display absorption components which show maximum depth at phase opposite primary eclipse (\\citealt{szkody90}).\\par  Currently, none of the proposed models have succesfully explained all the SW Sex peculiarities. The first model was proposed by \\citet*{honeycutt86}, who suggested the presence of an accretion disk wind. \\citet{williams89} invoked magnetic accretion, with single-peaked lines arising in accretion columns above the orbital plane, while \\citet{hellier94} introduced a gas stream which overflows the disk. \\citet{casares96} proposed an scenario in which matter is accreted on to the magnetic poles of a synchronously rotating white dwarf. More recently, \\citet{horne99} introduced the idea of a disk-anchored magnetic propeller. The latest model was proposed by \\citet{groot99}, who invoked the presence of a strong shock that occurs along the accretion stream trajectory, but inside the disk.\\par  LS Peg (S193) is a non-eclipsing SW Sex system with an orbital period of $P_{orb}\\simeq$ 4.19 h (\\citealt*{rodriguez98,martinez99,taylor99}). \\citet{stockman92} did not detect a significant level of circular polarization in the 5000--6000 \\AA~band after averaging only two 8 minute exposures. The reported polarization level was $-0.01 \\pm 0.05$ \\%.\\par ",
        "conclusions": "We have detected a clear EW modulation in the blue wing of H$_\\beta$, with a frequency of $\\nu_{EW}=43.0 \\pm 3.0$ d$^{-1}$. Similar flaring behavior has been observed in H$_\\alpha$, H$_\\beta$, and \\ion{He}{2} $\\lambda$4686 in the spectra of the SW Sex-type old nova BT Mon \\citep{smith98}, with a period of $\\simeq$~32 minutes, and more recently in the SW Sex star DW UMa \\citep{smith00}, where a 30-minute period is found. This flaring is typical of intermediate polar CVs (IPs, e.g. FO Aqr, \\citealt{marsh96}; see also \\citealt{hellier99}).  We have also discovered variable circular polarization in LS Pegasi. It is the first time that circular polarization is detected in an SW Sex system. Since the measured polarization comes from the continuum, we tentatively attribute its origin to cyclotron emission (see \\S~4.2 for discussion). We associate the period $P_p$ to the spin period of the white dwarf. This result seems to suggest an ``IP-like'' structure in LS Peg. Of the known IPs, only three show circular polarization: BG CMi \\citep{penning86}, PQ Gem (\\citealt{piirola93}; \\citealt{rosen93}; \\citealt{potter97}) and RX J1712.6-2414 \\citep{buckley95}. The reason why we only observe circular polarization in these IPs is that they harbor the most magnetic white dwarves ($B_1\\sim$ 2--8 MG, \\citealt{wickra91}). The frequency of the EW modulation ($\\nu_{EW}$) is slightly smaller than the assumed spin frequency ($\\nu_p$). The difference between them is $\\nu_p - \\nu_{EW}=5.7$ d$^{-1}$ which coincides with the orbital frequency. Hence, we identify $P_{EW}$ with the beat period (i.e. synodic) between the white dwarf spin period and the orbital period.  \\subsection{Accretion scenario}  \\citet{groot99} reports the observation of a bright spot in SW Sex itself (when the system was in low state) located well inside the outer disk radius, between 0.4--0.6 $R_{L_1}$, where $R_{L_1}$ is the distance from the inner Lagrangian point ($L_1$) to the white dwarf. He suggested that this hot region is the consequence of a strong shock between the gas stream and the disk, but he also pointed out that the mechanism producing the shock is not clear. From the results we present for LS Peg, we propose that this inner bright spot could be produced by the shock of the disk-overflown gas stream against the white dwarf's magnetosphere. The shock should occur close to the corotation radius (i.e. the distance from the primary at which the Keplerian and white dwarf angular velocities match, $r_{co}$) and above the disk plane. Current theories of accretion disk-magnetic field interaction require the magnetospheric radius to be nearly equal to the corotation radius (see e.g. \\citealt{ghosh78}).\\par Let's assume that the shock takes place close to $r_{co}$. Then, $r_{shock} \\approx r_{co}=f R_{L_1}$, where $f \\le 1$. The distance (expressed in units of the orbital separation, $a$) between the white dwarf and the $L_1$ point is given by $R_{L_1}/a=0.46(1+q)^{-1/3}$ \\citep{warner76}, while the corotation radius is: \\begin{equation} r_{co}=\\left[\\frac{G M_1 P^2_1}{4 \\pi^2}\\right]^{\\frac{1}{3}}, \\label{eq1} \\end{equation} where $M_1$ is the mass of the primary and $P_1$ its spin period. From Kepler's Third Law and assuming $r_{shock} \\approx r_{co}$, we obtain: \\begin{equation} P_1 \\approx 0.31 f^{3/2} P_{orb} \\label{eq2} \\end{equation} For LS Peg, $P_{orb}=0.174774$ d and $P_1=29.6$ min (from the circular polarization modulation), so $f \\approx 0.52$. This result is in good agreement with the derived bright spot location in SW Sex, suggesting that the magnetic shock mechanism could be also present in this system. Furthermore, we have detected variable $B$-band circular polarization in SW Sex itself (P. Rodr\\'\\i guez-Gil et al., in preparation) with a period of about 28 minutes. Entering this value together with the orbital period in (\\ref{eq2}), we get $f \\approx 0.60$, which is consistent with Groot's results.  \\subsection{The magnetic field}  The magnetospheric radius in IPs, $r_0$, is usually given by $r_0 \\approx 0.52 r_\\mu$ \\citep{ghosh79}, where $r_\\mu$ is the Alfv\\'en radius. Then, \\begin{equation} r_0 \\approx 0.52 r_\\mu = 4.1 \\mu^{4/7} {M_1}^{-1/7} {\\dot{M}}^{-2/7}, \\label{eq3} \\end{equation} where $\\mu$ is the magnetic moment of the primary, $M_1$ its mass, and $\\dot{M}$ the mass transfer rate. As we mentioned above, the magnetospheric radius should be comparable to the corotation radius, so let's assume that $r_0 \\approx r_{co}$. From (\\ref{eq1}) and (\\ref{eq3}), we obtain: \\begin{equation} \\mu \\approx 6.49 \\times 10^{-7} {M_1}^{5/6} {\\dot{M}}^{1/2} {P_1}^{7/6}. \\label{eq4} \\end{equation} Entering $P_1\\simeq1800$ s, and $M_1=0.75~\\mathrm{M}_\\odot$ (a plausible mass value for SW Sex systems), we finally get: \\begin{equation} \\mu \\approx 1.8 \\times 10^{25} {\\dot{M}}^{1/2}. \\label{eq5} \\end{equation} To derive the magnetic field of the white dwarf in LS Peg, we must know its radius ($R_1$), since $B_1=\\mu/{R_1}^3$. We can give an estimate of $R_1$ using the mass-radius relation for white dwarfs with carbon cores of \\citet{hamada61}, obtaining $R_1 \\simeq 7.2 \\times 10^8$ cm for a $M_1=0.75~\\mathrm{M}_\\odot$ white dwarf. On the other hand, we must enter a value for the mass accretion rate, $\\dot{M}$. From the results of \\citet{echevarria94}, we can consider that $\\dot{M} \\approx 10^{16}-10^{17}$ g s$^{-1}$ in SW Sex systems. Hence, the magnetic field intensity of the white dwarf in LS Peg should be $$B_1 \\approx 5-15~~\\mathrm{MG}.$$ This range of values is consistent with what is found in the IPs which show circular polarization. Nevertheless, we have to stress the fact that we see modulated circular polarization in the $B$-band. In PQ Gem ($B_1 \\approx 9$--21 MG, \\citealt{potter97}), little circular polarization is detected in that band. As cyclotron models predict, PQ Gem shows a level of circular polarization increasing redwards from the $V$-band, so it is not completely clear that the circular polarization we observe in LS Peg is related to continuum cyclotron emission. On the other hand, we can not completely rule out this possibility due to the fact that the polarization level in LS Peg is very low. It would be crucial to perform longer wavelength circular polarimetry to determine whether the circular polarization we see is of cyclotron emission origin.\\par We finally propose that magnetic accretion plays a fundamental role in the behavior of SW Sex stars, and that they are IP systems with mass accretion rates probably larger than in the rest of IPs in quiescence."
    },
    "0011/astro-ph0011183_arXiv.txt": {
        "abstract": " ",
        "introduction": "Binary population synthesis is a useful tool for studying  the statistical properties of stars, including the compact objects. In this  paper we wish to address the following questions: what is the distribution of masses of the compact objects? what is the relation between the stellar initial mass and the final mass of a compact object, when binary evolution is taken into account? given the distribution of compact object masses, what are the relative numbers of different types of objects (neutron stars, quark stars, black holes)?  In section 2 we shortly describe the population synthesis code used here, in section 2 we present the results, and finally in section 4 we present the conclusions.    \\begin{figure*} \\includegraphics[angle=-90,width=0.95\\textwidth]{diagram1.ps} \\caption{The various routes of compact object formation (compact object = a neutron star, a strange star or a black hole). } \\end{figure*} ",
        "conclusions": ""
    },
    "0011/astro-ph0011460_arXiv.txt": {
        "abstract": "We have used the Very Long Baseline Array (VLBA) and the Very Large Array (VLA) to image Galactic neutral hydrogen in absorption towards four compact extragalactic radio sources with 10 milliarcsecond resolution.  Previous VLBA data by Faison et al (1998) have shown the existence of prominent structures in the direction of the extragalactic source 3C~138 with scale sizes of 10-20 AU with changes in HI optical depth in excess of 0.8 $\\pm$ 0.1.  In this paper we confirm the small scale \\hi\\ optical depth variations toward 3C~147 suggested earlier at a level up to 20 \\% $\\pm$ 5\\% .  The sources 3C~119, 2352+495 and 0831+557 show no significant change in \\hi\\ optical depth across the sources with one sigma limits of 30\\%, 50\\%, and 100\\%.  Of the seven sources recently investigated with the VLBA and VLA , only 3C~138 and 3C~147 show statistically significant variations in HI opacities.  Deshpande (2000) have attempted to explain the observed small-scale structure as an extension of the observed power spectrum of structure on parsec size scales.  The predictions of Deshpande (2000) are consistent with the VLBA \\hi\\ data observed in the directions of a number of sources, including 3C~147, but are not consistent with our previous observations of the \\hi\\ opacity structure toward 3C~138. ",
        "introduction": "Structure in the interstellar medium (ISM) is observed on many different size scales, from large kiloparsec scale superbubbles (Dewdney and Lozinskaya 1994) to electron density fluctuations in the ionized neutral medium only a few hundred kilometers in size (Armstrong, Rickett, and Spangler 1995).  Previous studies have attempted to characterize the structure of the neutral component of the ISM by observing neutral hydrogen in emission over a range of angular scales (e.g. Green 1993; Crovisier and Dickey 1983).  These studies have been limited by the angular resolution of single dish telescopes and connected-element interferometers to size scales from about 0.1 parsec to several hundred parsecs.  The results of these studies suggested that the neutral ISM is smooth on scales smaller than about 0.1 pc.  Due to the typically low brightness temperature of \\hi\\ in emission on small scales, \\hi\\ 21 cm emission can not be detected at milliarcsecond scales using Very Long Baseline Interferometry (VLBI) techniques.  However, it is possible to image Galactic \\hi\\ in absorption towards bright extragalactic continuum sources, and thereby study the structure of the neutral ISM on angular size scales from the resolution of the VLBI array (about 5 mas at 21 cm for the VLBA) up to the angular size of the continuum source ($\\propto$ 50 to 200 mas for the sources in this paper).  At typical distances to the absorbing gas of 500 to 1000 pc, these angular scales correspond to tens to hundreds of Astronomical Units.  The first VLBI observations of very small scale structure in the neutral ISM were published by Dieter \\etal\\ (1976).  Observing the source 3C147 with a single baseline from Hat Creek to the Owens Valley Radio Observatory (150 km), they observed variations in the ratio of the \\hi\\ absorption line visibility amplitude to the 21 cm continuum visibility amplitude as a function of projected baseline.  This change is caused by non-uniform coverage of the \\hi\\ absorption across the face of the source (Radhakrishnan \\etal\\ 1972).  Based on the galactic latitude of the source and a model of the continuum emission components, they inferred that there was structure in the absorbing \\hi\\ with a physical size of 70 AU.  These results, based on modeling of sparse visibility data, were qualitative, but suggested that there is structure on very small scales in the diffuse neutral ISM.  Subsequent VLBI observations of AU scale \\hi\\ structure were published by Diamond \\etal\\ (1989).  They observed the three extragalactic radio continuum sources 3C~147, 3C~138 and 3C~380 using three large telescopes of the European VLBI Network.  The visibility data from the observations were not sufficient to image the sources and the \\hi\\ absorption distribution.  However, the ratio of the line to continuum visibility amplitudes confirmed the non-uniform \\hi\\ absorption towards 3C~147 observed by Dieter \\etal\\ (1976). Large variations towards 3C~138 and moderate variations towards 3C~380 were discovered. Diamond \\etal\\ inferred structure in the \\hi\\ gas of 25 AU with densities of $5 - 10 \\times 10^4$ \\pcm, assuming that the line of sight dimension is equal to the measured transverse dimension.  These conclusions were based on modeling the visibility data rather than directly imaging the \\hi\\ optical depth structure.  Confirmation of AU scale structure in Galactic \\hi\\ independent of the VLBI data was published by Frail \\etal\\ (1994).  They observed average variations in the total \\hi\\ column density of 10-15\\% towards nine pulsars at three observing epochs over a period of 1.7 years.  Each pulsar in the Frail \\etal\\ study probed a range of physical size scales transverse to the line of sight from 5 to 100 AU. The fact that column density variations over time were observed towards all of the pulsars suggests that AU scale structure is common throughout the ISM. Also, the fact that Frail \\etal\\ observed changes in the total \\hi\\ column density but not in the overall line shape suggests that the AU scale structure is kinematicly related to the diffuse cold neutral medium as a whole.  The first images of AU scale \\hi\\ structure in absorption were published by Davis, Diamond, and Goss (1996) (Henceforth DDG).  They used the MERLIN array with five EVN telescopes to image 3C~147 and 3C~138 in \\hi\\ optical depth at an angular resolution of 50 $-$ 150 mas.  They observed significant optical depth variations in one of the two \\hi\\ absorption lines towards both objects.  The largest optical depth variation of $\\Delta\\tau_{HI} = 0.16 \\pm 0.03$ across the source was observed in the 1.3 \\kps\\ absorption line towards 3C~138.  A prominent optical depth change of $\\Delta\\tau_{HI} = 0.09 \\pm 0.01$ was also observed in the -8 \\kps\\ absorption line towards 3C~147.  There is also evidence for AU scale structure in the diffuse molecular gas.  Moore and Marscher (1995; as well as Marscher \\etal\\ 1993) have discussed changes observed in the column density of H$_2$CO towards the compact objects 3C~111, BL~Lac, and NRAO~150.  They observed the three sources over 3.4 years with the VLA in the 6 cm H$_2$CO line. The motion of the sources due to parallax as well as the proper motions of the absorbing gas caused the relative line of sight through the molecular gas to the extragalactic sources to change with time. They observed significant variations in the H$_2$CO opacities towards 3C~111 and NRAO~150, and inferred structure on linear scales $\\sim$ 10 AU and H$_2$ densities $\\sim 10^6$ cm$^{-3}$, sizes and densities similar to those observed for the small scale \\hi\\ structure.  Heiles (1997) has proposed a geometrical model to explain the large variations in column density.  The geometrical model assumes that the structures are due to filaments or sheets observed edge-on, causing larger variations in column densities over smaller scales to be observed and implying much higher volume densities than are actually present.  If the observed structure arises from edge-on sheets, the optical depth images should statistically show smaller-scale variations along the axis perpendicular to the sheet compared to variations along the axis parallel with the sheet.  Faison, Goss, Diamond, and Taylor (1998), hereafter FGDT, published high angular resolution (10-20 mas) VLBA images of \\hi\\ absorption towards 3C~138, 2255+416, and 0404+768.  Pronounced \\hi\\ optical depth variations were confirmed towards 3C~138, consistent with Diamond \\etal\\ (1989) and DDG.  Moderate \\hi\\ variations were observed towards the small angular size source 2255+417, and no significant variations were observed towards 0404+768, suggesting a minimum size of the \\hi\\ structure of a few tens of AU.  Deshpande, Dwarakanath, and Goss (2000) have evaluated the structure observed in \\hi\\ absorption towards Cas A over a linear scale of 4pc to 0.02 pc. Deshpande (2000) has taken the slope of the power spectrum of fluctuations of $\\alpha = -2.75$ and extrapolated to scales smaller by factors of three orders of magnitude (sizes scales smaller than 100 AU) and predicts that optical depth variations as large as 0.1 should be expected on scales of tens of AU. The cause of these fluctuations is due to gradients in the distribution caused by the ``red'' spectrum (larger scales have more power) of the power spectrum of the HI. The fluctuations observed by the VLBA in the direction of 3C~138 are almost an order of magnitude larger than this prediction. In addition, two dimensional HI features are observed.  In this paper we present \\hi\\ optical depth images towards 3C~147, 3C~119, 2352+495, and 0831+557, and we discuss our current knowledge of AU scale structure in the diffuse neutral ISM. ",
        "conclusions": "The optical depth variations towards 3C~147 are shown in the position-velocity plot for 3C~147, shown in Figure~\\ref{3c147lv}.  The vertical axis of this position-velocity corresponds to the axis along the source from the NW to the SE as indicated in Figures~\\ref{3c147taumap119} and \\ref{3c147taumap140}.  As can be noticed in the individual optical depth images, the optical depth variations along the axis of the source are larger in \\hi\\ absorption line near $-8.3$ than in the line near $0.4$ \\kps.  The fluctuations in the opacity in Figure~\\ref{3c147lv} far away from the absorption lines are a good indication of the noise in the optical depth at a given position along the axis of the source.  This position-velocity plot can be compared with the position-velocity plot for 3C~138 from FGDT.  Although there are obvious optical depth variations in each absorption line towards 3C~147, the striking optical depth structure changes in velocity observed towards 3C~138 are not observed in the 3C~147 position-velocity image.  In the 3C~138 position-velocity plot, the velocity features that slope diagonally across are quite striking and indicate that we are observing real, three dimensional features coherent in velocity.  The constraints on the line of sight magnetic field towards these sources suggest that the magnetic field towards these sources is not enhanced relative to the average interstellar magnetic field of $40 - 100$ \\uG .  Diamond \\etal\\ (1989) determined that field strengths of 200 \\uG\\ would be required to confine the small scale structure observed towards 3C~138 relative to the pressure of the diffuse ISM. The upper limits on the line of sight magnetic field towards 3C~147 and 3C~138 are an order of magnitude below this value.  In conclusion, we have imaged Galactic \\hi\\ in absorption towards four extragalactic continuum objects at 10 milliarcsecond resolution.  The -8.3 \\kps\\ \\hi\\ absorption line towards the source 3C~147 shows variations of $10$ to $20$ \\% in \\hi\\ column density across the source ($1 \\sigma$ error of 5\\%), and variations in optical depth from about 0.82 to 1.1 (at about a $4 \\sigma$ level).  No significant variations \\hi\\ optical depth are observed across the source 3C~119 to within a $1 \\sigma$ error of 0.2 in optical depth and a 10\\% change in \\hi\\ column density.  The source 2352+495 shows no significant \\hi\\ column density variations in column density to within a $1 \\sigma$ limit of 30\\%.  The high Galactic latitude source 0831+557 has weak \\hi\\ absorption and shows no significant column density variations to within a $1 \\sigma$ limit of 50\\%.  All of these results are consistent with the prediction of Deshpande (2000) that optical depth variations as large as $\\delta\\tau \\approx 0.1$ on scales of 100's of AU.  However, the source 3C~138 is unique among the objects we have observed in showing very large \\hi\\ optical depth variations of $\\Delta\\tau \\approx 0.1$ on scales of 10's of AU.  It is of considerable interest to probe the cold, neutral medium on intermediate scales between the few tens of AU probed by VLBI observations and the 0.1 pc scales probed by connected-element interferometers such as the WSRT and the VLA.  The only observations published to date on these angular scales between the resolution of the VLBA at about $10-500$ mas and the VLA at about $1-30$ \\arcsec\\ are observations of ISM optical lines towards binary stars (e.g., Lauroesch \\etal\\ 1999; Watson and Meyer 1996) separated by 500 to 30,000 AU.  However, most of these observations are of variations in \\ion{Na}{1} column density, which may not reliably trace the \\hi\\ column density.  Observations of \\hi\\ column density variations should be carried out on these size scales to compare the AU scale structure with the larger scale structure.  These observations could be carried out with the MERLIN array and also by observing \\hi\\ column density variations towards pulsars as done by Frail \\etal\\ (1994).  The Green Bank Telescope will be an especially sensitive instrument for observing \\hi\\ in absorption towards pulsars.  Also, more pulsar and VLBI studies of AU scale structure should be carried out in the southern hemisphere, outside the field view of the VLBA.  To date no observations of AU scale structure in \\hi\\ have been made from $l=187$\\dg\\ to $l=20$\\dg."
    },
    "0011/astro-ph0011526_arXiv.txt": {
        "abstract": "We review current astrophysical bounds on MeV sterile neutrinos, and then we discuss why a sterile keV neutrino is a natural warm dark matter candidate. ",
        "introduction": "The experimental evidence for neutrino masses and mixing is overwhelming~\\cite{ellis}, and if all the present day experiments are correct, then there exists at least one sterile neutrino besides the normal 3 active ones. The mass differences are found to be in the sub-eV range, but one could naturally imagine additional sterile neutrinos, and since both masses and mixing angles essentially are free parameters, one naturally asks the question, which ones are already excluded - or are there indications that heavy sterile neutrinos exist?  Let us first review the current bounds on MeV sterile neutrinos, and then later discuss how a keV sterile neutrino naturally could be produced in the early universe in just the right amount to be a warm dark matter candidate. ",
        "conclusions": ""
    },
    "0011/astro-ph0011076_arXiv.txt": {
        "abstract": "In previous studies we have shown that the optical variability of quasars increases, on average, with redshift. We explained this dependance in terms of a hardening of the spectrum during bright phases, coupled with the increase of the rest-frame frequency for increasing redshift. We re-analize now these correlations on the basis of new light curves of PG quasars, recently  published by the Wise Observatory group. ",
        "introduction": "Although variability plays a key role in constraining the models of the central engine of QSOs, little is known about the physical origin of luminosity variations.  The most diverse variability mechanisms have been proposed in the past, including supernovae explosions (Aretxaga et al. 1997), instabilities in the accretion disk (Kawaguchi et al. 1998), and gravitational lensing due to intervening matter (Hawkins 1993). So far, most of the information about the characteristics of variability derives from the statistical analysis of single band light curves of magnitude limited QSO samples. The correlation of variability with either intrinsic luminosity ($v-L$) or redshift ($v-z$) is affected by the strong correlation between luminosity and redshift ($L-z$), present in these samples. The results of these analyses also depend on the specific variability index adopted, as shown by Giallongo et al. (1991), who found a positive v-z correlation through a variability index based on the rest-frame structure function, later confirmed by Cristiani et al. (1996). The increase of variability with the rest-frame observing frequency found by Di Clemente et al. (1996), supports the suggestion by Giallongo et al. (1991) that QSOs at high redshift appear more variable, on average, since they are observed in a higher rest frame frequency. The dependence of variability on frequency is associated with the hardening of the spectral energy distribution (SED) during the bright phases (Cutri et al. 1985, Kinney et al. 1991, Edelson et al. 1990). Tr\\`evese et al. (1999) have shown that a hardening of the SED in the bright phase occurs, on average, in the statistical, magnitude limited, sample of AGNs of the SA 57 (Tr\\`evese et al. 1989, 1994, Bershady et al. 1998). They also found that the slope a and its variations associated with the luminosity changes, are consistent on average with temperature changes of a black body.  \\subsection{Data and analysis}  We present preliminary results of a new analysis of the spectral variability of QSOs in the optical band. It is based on a data set made available by Giveon et al. 1999, consisting of the light curves of a sample of 42 nearby and bright QSOs ($z<0.4$, $B<16$ mag) belonging to the Palomar-Green (PG) sample, which were monitored for 7 years, with a typical sampling interval of 39 days, in the Johnson Cousins B and R bands, with the 1 m telescope of the Wise Observatory. Here the quantities adopted for the analysis are defined. The results are discussed in Section 2, together with the relevant figures.  \\subsection{Variability index}  We define the structure function as (see Di Clemente et al. 1996) $S_1=[\\pi/2(\\langle|m(t+\\tau)-m(t)|\\rangle^2-\\sigma_n^2)]^{1/2}$, where $m(t)$ is either the $B$ or the $R$ magnitude, $t$ is the rest-frame time, $\\tau$ is the time lag between the observations, $\\sigma_n$ is the relevant r.m.s. noise and the angular brackets indicate the average taken over all the pairs of observations separated by a time interval $\\tau\\pm\\Delta\\tau$.  The value of $\\Delta\\tau$ is the result of a trade-off between time resolution and statistical uncertainty.  In the following we define, for each QSO, two {\\bf variability indexes} $v_B$ and $v_R$ equal to the relevant structure function in the $B$ and $R$ bands, computed for $\\tau=2\\pm 0.6$ yr. The average values of $v_R$ and $v_B$ over the ensemble of 42 QSOs, versus the relevant average rest-frame frequency, are reported in Figure 1, which is adapted from Di Clemente et al. (1996).  \\subsection{Instantaneous spectral slope}  For each QSO we compute the {\\bf instantaneous spectral slope} $\\alpha(t)\\equiv\\log(f_{\\nu_B}/f_{\\nu_R})/\\log(\\nu_B/\\nu_R)$, regarding as synchronous the observations within a time lag of 9 hours. Emission line corrections are not considered in the present preliminary analysis.  \\subsection{Spectral variability parameter}  Considering all possible pairs of observations, separated by a time interval $\\tau_{ij}\\equiv\\tau_i-\\tau_j$, we define a spectral variability indicator: $\\beta(\\tau)\\equiv[\\alpha(t+\\tau)-\\alpha(t)]/[\\log L_B(t+\\tau)-\\log L_B(t)]$ where $L_B(t)$ is the absolute luminosity in the blue band. In the present analysis we consider, for each object, the {\\bf spectral variability parameter} $\\beta_m$ defined as the average of all the $\\beta(\\tau)$ values for $\\tau\\le 1000$ days. ",
        "conclusions": "The analysis of the spectral variability of the present statistical subsample of PG QSOs shows that: \\begin{itemize} \\item the amplitude of variability increases with the rest-frame frequency; \\item a (positive) $v-z$ correlation  appears in the present sample,  despite the (negative) $v-L$ and the (positive) $L-z$ correlations; \\item the increase of variability with frequency is consistent with the trend previously found by Di Clemente et al. (1996) on the basis of different QSO samples and confirms the explanation of the $v-z$ correlation suggested by Giallongo et al. (1991); \\item the spectral slope increases on average in the bright phases (intra-QSO $\\alpha-L$ correlation); \\item the spectral slope is steeper, on average, for fainter QSOs (inter-QSO $\\alpha-L$ correlation); \\item the two latter points suggest that the gas temperature is  responsible for both the effects; \\item temperature variations of a black body with an average temperature of the order of $2.5\\cdot 10^4$ K are consistent for the intra-QSO $\\alpha(t) - L_B(t)$ correlation. \\item different average black body temperatures may account for the inter-QSO $\\alpha - L_B$ correlation. \\end{itemize}"
    },
    "0011/astro-ph0011289_arXiv.txt": {
        "abstract": "Calculations of the statistical equilibrium of magnesium in the solar photosphere have shown that NLTE populations hardly affect Mg line formation in the Sun. However, in metal-poor dwarfs and giants the influence of electron collisions is reduced, and the ultraviolet radiation field, enhanced due to reduced background line opacity, results in more pronounced NLTE effects. In the photosphere of a cool star excitation and ionization due to collisions with neutral hydrogen can outweigh electron collisions.  Analyses based on NLTE populations lead to significantly higher Mg abundances than those calculated from LTE. We calculate magnesium abundances in 10 cool dwarfs and subgiants with metallicities from $-2.29$ to $0.0$. The results are based on spectra of high-resolution and high signal-to-noise ratio. Stellar effective temperatures are derived from Balmer line profiles, surface gravities from Hipparcos parallaxes and the wings of the \\ion{Mg}{i}b triplet, and metal abundances and microturbulence velocities are obtained from LTE analyses of \\ion{Fe}{ii} line profiles. For stars with metallicities between $-2.0 <$ [Fe/H] $< -1.0$ abundance corrections $\\Delta{\\rm [Mg/H]_{NLTE-LTE}} \\sim 0.05 - 0.11$ are found. As expected the corrections increase with decreasing metal abundance, and they increase slightly with decreasing surface gravity. We also calculate the  statistical equilibrium of magnesium for series of model atmospheres with different stellar parameters and find that $\\Delta{\\rm [Mg/H]_{NLTE-LTE}}$ increases with effective temperature between 5200 and 6500 K. For extremely metal-poor stars the abundance corrections approach $\\Delta{\\rm [Mg/H]_{NLTE-LTE}} \\sim 0.23$ at [Fe/H] $\\sim -3.0$. ",
        "introduction": "Abundance patterns of $\\alpha$-elements, such as Mg, Si and Ca, in stars with different metallicities are important for understanding the chemical evolution of Galaxy. The first stellar generations are supposed to have produced mostly $\\alpha$-elements during massive supernova explosions of type II. Studies of $\\alpha$-element synthesis in SNe~II have been undertaken by Arnett (\\cite{ARNETT91}), Thielemann et al. (\\cite{THIELEMANN96}), Woosley \\& Weaver (\\cite{WOOSLEY95}), and Nakamura et al. (\\cite{NAKA99}). Due to small differences in the way stellar winds and semi-convection are treated and in the specification of the {\\em mass cut} and explosion mechanism their predictions differ, mainly in the respective Fe yields but also in the pre-explosion yields of Mg. There is no question that Mg, in principle, is less affected by the fine-tuning of SN~II explosions. Therefore, inasmuch as the products of SN~II nucleosynthesis are mixed into interstellar space \\Mag\\ should constitute a reliable reference of the early evolution of the Galaxy. Due to the number of strong absorption lines found in the spectra of even the metal-poor stars, \\ion{Mg}{i} is easier to observe than e.g. \\ion{O}{i}. However, it shares the disadvantage of most neutral metals in the atmospheres of moderately cool stars with \\ion{Mg}{ii} being the dominant ionization stage at temperatures above 5000 K. Consequently, neutral magnesium is sensitive to deviations from {\\em local thermodynamic equilibrium}, particularly as its ionization balance is dominated by large photoionization cross-sections from the \\Tripp\\ state. In metal-poor stars reduced line-blanketing produces an increased UV radiation field that, in combination with the large cross-sections, leads to strong photoionization rates; this makes \\ion{Mg}{i} even more sensitive to NLTE and affects any careful abundance analysis of Mg in these objects. Moreover, the density of free electrons decreases roughly in proportion to metal abundance, and the collision rates become correspondingly smaller. Together this change of microscopic interaction processes should lead to substantially stronger deviations from LTE in metal-poor stars at optical depths $\\log$\\,\\tauc between $-3$ and $0$, if the reduced electron collisional interaction is not compensated by hydrogen collisions (Baum\\\"uller \\& Gehren \\cite{BAUM96},\\cite{BAUM97}).  In our present study, we use several \\ion{Mg}{i} lines chosen from the solar spectrum which were carefully analyzed previously (Zhao et al. \\cite{ZHAO98}, hereafter Paper~I). The standard atmospheric models we use in our present analysis are described in Sect. 2. We start with test calculations including the standard model atom of Paper I, and we fit the calculated NLTE equivalent widths for models of different stellar parameters. We determine Mg {\\em abundance corrections} due to differences between NLTE and LTE analyses for series of model atmospheres with the different stellar parameters in Sect. 3. All \\ion{Mg}{i} lines considered here are reproduced using standard NLTE line formation techniques with the radiative transfer solved in the Auer-Heasley scheme, taking into account the interaction processes between all levels of the Mg model atom. A previous investigation (Paper~I) has demonstrated how the line profiles observed in the solar spectrum can be used to improve the atomic model of neutral magnesium. This model is the subject of further scrutiny in the light of information coming from comparison of test calculations with line profiles in the spectra of metal-poor stars. We then calculate magnesium abundances of 10 cool dwarfs and subgiants with metallicities from $-2.29$ to $0.0$ using both LTE and NLTE line formation techniques in Sect. 4. Our conclusions are given in last section. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1_1.ps}}\\\\ % \\resizebox{\\hsize}{!}{\\includegraphics{fig1_2.ps}}   % \\caption[]{NLTE corrections of the magnesium abundances versus [Fe/H] for the model grids, where $\\Delta$[Mg/H] = [Mg/H]$_{\\rm NLTE}$ -- [Mg/H]$_{\\rm LTE}$. Top: \\Teff = 5780 K, \\logg = 3.50 and 4.44. Bottom: \\logg = 4.50 (4.44 for \\Teff = 5780 K), \\Teff = 5200, 5780, and 6500 K.} \\label{nlte_corr} \\end{figure} \\begin{table*} \\caption[]{Abundance differences between NLTE and LTE obtained by fitting calculated LTE equivalent widths of \\ion{Mg}{i} lines using the same parameters for various stellar model atmospheres based on our line-blanketed model grid. Results refer to $\\Delta$[Mg/H] = [Mg/H]$_{\\rm NLTE}$ - [Mg/H]$_{\\rm LTE}$} \\label{nltediff} \\begin{tabular}{lrrrrrrrrrrrr} \\hline \\noalign{\\smallskip} \\Teff& \\logg &[Fe/H]&4571\\AA&4703\\AA&4730\\AA&5528\\AA&5711\\AA&7657\\AA&8806\\AA&8213\\AA&11828\\AA&$\\overline{\\Delta[Mg/H]}\\pm\\sigma$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 5200 &  4.50 &  0.0 & 0.022 & 0.014 & 0.021 & 0.010 & 0.015 & 0.035 & 0.001 & 0.042 & 0.004  & 0.018$\\pm$0.014\\\\ 5200 &  4.50 & -1.0 & 0.060 & 0.038 & 0.051 & 0.022 & 0.045 & 0.065 & 0.019 & 0.115 & 0.025  & 0.049$\\pm$0.030\\\\ 5200 &  4.50 & -2.0 & 0.091 & 0.039 & 0.058 & 0.028 & 0.052 & 0.055 & 0.045 & 0.098 & 0.045  & 0.057$\\pm$0.023\\\\ 5200 &  4.50 & -3.0 & 0.113 & 0.062 & 0.068 & 0.058 & 0.060 & 0.050 & 0.075 & 0.095 & 0.075  & 0.073$\\pm$0.020\\\\[2mm]  5780 &  4.44 &  0.0 & 0.047 & 0.028 & 0.042 & 0.023 & 0.050 & 0.060 & 0.004 & 0.076 & 0.019  & 0.039$\\pm$0.022\\\\ 5780 &  4.44 & -1.0 & 0.087 & 0.073 & 0.088 & 0.060 & 0.104 & 0.130 & 0.041 & 0.202 & 0.050  & 0.093$\\pm$0.049\\\\ 5780 &  4.44 & -2.0 & 0.128 & 0.123 & 0.111 & 0.099 & 0.117 & 0.120 & 0.079 & 0.234 & 0.090  & 0.122$\\pm$0.045\\\\ 5780 &  4.44 & -3.0 & 0.182 & 0.190 & 0.172 & 0.180 & 0.172 & 0.200 & 0.163 & 0.430 & 0.157  & 0.205$\\pm$0.085\\\\[2mm]  5780 &  3.50 &  0.0 & 0.060 & 0.033 & 0.062 & 0.028 & 0.090 & 0.115 & 0.016 & 0.095 & 0.026  & 0.058$\\pm$0.035\\\\ 5780 &  3.50 & -1.0 & 0.101 & 0.081 & 0.120 & 0.054 & 0.151 & 0.180 & 0.018 & 0.252 & 0.053  & 0.112$\\pm$0.073\\\\ 5780 &  3.50 & -2.0 & 0.147 & 0.157 & 0.137 & 0.118 & 0.136 & 0.120 & 0.045 & 0.264 & 0.085  & 0.134$\\pm$0.060\\\\ 5780 &  3.50 & -3.0 & 0.201 & 0.224 & 0.206 & 0.211 & 0.207 & 0.210 & 0.212 & 0.429 & 0.210  & 0.234$\\pm$0.073\\\\[2mm]  6500 &  4.50 &  0.0 & 0.042 & 0.041 & 0.046 & 0.045 & 0.083 & 0.100 & 0.042 & 0.095 & 0.047  & 0.060$\\pm$0.025\\\\ 6500 &  4.50 & -1.0 & 0.077 & 0.102 & 0.088 & 0.090 & 0.123 & 0.150 & 0.051 & 0.188 & 0.072  & 0.105$\\pm$0.043\\\\ 6500 &  4.50 & -2.0 & 0.096 & 0.142 & 0.101 & 0.120 & 0.105 & 0.120 & 0.075 & 0.178 & 0.094  & 0.115$\\pm$0.031\\\\ 6500 &  4.50 & -3.0 & 0.148 & 0.178 & 0.162 & 0.182 & 0.179 & 0.200 & 0.178 & 0.302 & 0.175  & 0.189$\\pm$0.045\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "The variation of [Mg/Fe] with the stellar mean metallicity [Fe/H] contains information about the chemical evolution of the Milky Way. Previous investigations of magnesium based on LTE analyses by a number of researchers (\\eg Magain \\cite{MAGAIN87}; Hartmann \\& Gehren \\cite{HARTMANN88}; Fuhrmann et al. \\cite{FUHRMANN95}; Fuhrmann \\cite{F98}) show that magnesium has a solar abundance ratio in disk stars ([Fe/H] $\\geq$ $-$0.5), and then [Mg/Fe] becomes constant at +0.3 \\ldots 0.4 in metal-poor stars ([Fe/H] $\\leq  -0.5$). As shown by Fuhrmann (\\cite{F98}), this overabundance is a distinct result of population membership. Our new results based on the NLTE analysis of 10 cool unevolved stars demonstrate that the magnesium overabundance with respect to iron becomes even larger by roughly 0.1 dex for stars with reduced stellar metal abundance [Fe/H] due to photoionization dominating over collisions. The reduction of the surface gravity results in a decreased efficiency of collisions by electron and hydrogen atoms which also leads to small NLTE effects. The strongest departures from LTE are found for metal-poor stars with high temperature and low gravity, i.e. {\\em turnoff} stars.  Fortunately, the trend obtained from LTE abundance analyses of magnesium in cool stars (Fuhrmann \\cite{F98}) is confirmed. The significant [Mg/Fe] ratio in metal-poor stars is found from NLTE spectrum analyses, with the increased level indicating a shift to even higher stellar masses in the parent population (III). Such a trend was also found by Fuhrmann et al. (1995) or Audouze \\& Silk (1995) who investigated the initial mass function based on the $\\alpha$-element enrichment of population II. Recent investigations of the \\ion{Fe}{ii}/\\ion{Fe}{i} equilibrium (see Th\\'evenin \\& Idiart \\cite{TI99}), however, seem to demonstrate that [Mg/Fe] may be more affected by deviations from LTE when using \\ion{Fe}{i} linese to derive iron abundances. We recall that in our analysis we used only \\ion{Fe}{ii} linese which are insensitive to NLTE effects.  In order to confirm the present results for magnesium, and in particular the mean abundance differences between LTE and NLTE it is necessary to reinvestigate a larger sample of stars with various metallicities based on the NLTE analysis, especially for the metal-poor stars of the thick disk and halo. Such an investigation should, however, include an \\ion{Fe}{i} NLTE analysis. Then it will be an important step improving our present understanding of stellar nucleosynthesis and the chemical evolution of the Galaxy."
    },
    "0011/astro-ph0011240_arXiv.txt": {
        "abstract": " ",
        "introduction": "In recent years, new aspects of rotating neutron stars have been revealed in about 1000 pulsars. Eleven X-ray pulsars\\cite{bt99} and eight $\\gamma$-ray pulsars\\cite{th99} have been detected in the past several years. Among these new objects, some exhibit quite different behavior in their pulse periods.\\cite{obs} The measurement of the period and its time derivative yields evidence of ultra-magnetized stars, possibly representing magnetars.\\cite{dt} Motivated by the recent observational situation, theoretical models have been studied. As for high-energy pulsars, two general classes of models have been proposed. One is the polar cap model\\cite{pol} and the other is the outer gap model.\\cite{out} The main difference between these two models is in the assumed region of the acceleration of charged particles responsible for the radiation. Both models partially explain some observational features of the  $\\gamma$-rays. They will be discriminated after including more detailed radiation processes. Future observation may determine their validity.   An important element to be included in theoretical models is general relativistic effects, which are in particular crucial for polar cap models, since acceleration occurs under strong gravity near the surface of neutron stars. Gonthier and Harding\\cite{gh94} considered the effects on the magnetic field configuration only. Their concern is the curvature radiation and the attenuation of pair production in a strong magnetic field. These processes result in a pair cascade and explain some aspects of pulsar radiation, including high-energy pulses in the $\\gamma$-ray range. In addition to the magnetic fields, rotationally induced electric fields play an important role in the polar cap region (see, e.g., Ref.~\\citen{meszaros}). Charged particles are ripped off the surface and accelerated along the magnetic field lines by the electric fields. The magnetosphere is thereby eventually filled with charges. The accelerated particles may be seeds of subsequent curvature radiation. Muslimov and Tsygan\\cite{mt} discussed general relativistic effects not only in the case of magnetic fields but also electric fields. They derived general expressions including multipoles of arbitrary order using hypergeometric functions, assuming a vacuum outside the star. Their work, however, is limited to only analytic forms, and therefore not easy, e.g., to compare with the standard results in flat space-time. Order estimates of general relativistic effects are also lacking. In this paper, we derive analytic solutions again for both the electric and magnetic fields around a rotating star endowed with an aligned dipole magnetic field. The resultant expressions are rather cumbersome, and for this reason approximate expressions are also given. Such forms provide an estimate of the corrections to the results in flat space-time, as well as a concise, practical application. We also give detailed discussion concerning the difference between our results and those in Minkowski space-time. This discussion may become important in the future, with progress in observational technology.  As shown in Ref.~\\citen{kok2}, the deviation from spherical space-time is less than $ 10^{-3}$ if the rotation period is longer than 10 msec and the magnetic field at the surface is less than  $ 10^{16}$ gauss. Therefore, the electric and magnetic fields are determined by solving the Maxwell equations in a fixed background space-time. The appropriate space-time metric is that for an external field surrounding a slowly rotating star. We can neglect the second-order rotational effects, except in the case of rapidly rotating stars. We also restrict ourselves to a stationary configuration, that is, the case in which the magnetic dipole moment $\\mib{\\mu}$ is aligned with the angular velocity $\\mib{\\Omega}$. This leads to the following form $A_{\\mu} = (A_{t},0,0,A_{\\phi})$ for the four-potential (see Ref.~\\citen{bgsm} and references therein), where $A_{t}$ is related to the rotationally induced electric field, and therefore $A_{t} \\sim O \\left(  \\Omega \\right)\\times  A_{\\phi}  $. Detailed calculations to solve the Maxwell equations are given in \\S 2. Approximate expressions of these solutions are discussed in \\S 3. Implications of the general relativistic effects with regard to the acceleration of charged particles and radiation in vacuum gaps are investigated in \\S 4. Finally, we give discussion in \\S 5. Throughout the paper, we use units in which $c=G=1$. ",
        "conclusions": "Recent observations of compact stars have given remarkable results that demand the refinement of theoretical models. Inspired by this, we have reconsidered an exterior electromagnetic field surrounding a rotating star endowed with an aligned dipole magnetic field in the context of general relativity. The electromagnetic fields were derived in analytic and approximate forms. We found that the expressions calculated in the flat space-time are accurate within a factor of approximately 2. We have not calculated the emission and propagation of radiation from polar caps of pulsars, but rather have discussed the implications for the underlining physical processes. We have found that the general relativistic effects increase the strength of electric fields and decrease the curvature radii of the magnetic field lines. Both of these factors contribute to increase the photon energy emitted from charged particles. The magnitude of the correction is of order $ M/R. $ Another important general relativistic effect, which has not been considered here, is the redshift factor. The observed energy is shifted to a lower value by a factor of $ M/R. $ It is not clear whether or not all these general relativistic effects are canceled. It is important to construct detailed models of pulsar radiation, taking these factors into account.   Although we have restricted our investigation to a rotating star in a vacuum, it seems that actual neutron stars are surrounded by plasma. Hence, it is important to investigate the acceleration of charged particles and the radiation taking into account the plasma distributions around stars. A general relativistic analysis using a certain pulsar model which specifies the plasma distribution has been given by Muslimov and Tsygan.\\cite{mt2} The general relativistic effects in pulsar models are not yet clear, since the magnitude of the effects significantly depends on the plasma distribution. It is necessary to discuss the effects in a more general framework. This will be the subject of future investigation."
    },
    "0011/astro-ph0011130_arXiv.txt": {
        "abstract": "We show that models for screening of nuclear reactions in the Sun can be tested by means of helioseismology. As well known, solar models using the weak screening factors are in agreement with data. We find that the solar model calculated with the anti screening factors of Tsytovitch is not consistent with helioseismology, both for the sound speed profile and for the depth of the convective envelope. Moreover, the difference between the no-screening and weak screening model is significant in comparison with helioseismic uncertainty. In other words, the existence of screening can be proved by means of helioseismology. ",
        "introduction": "\\label{intro}  The solar neutrino problem is so important that any aspect of solar, plasma and nuclear physics pertinent to it has to be deeply investigated before definitive conclusions can be drawn. In this respect, screening of the charges of the reacting nuclei due to free charges in the solar plasma is of some interest.  The study of screened nuclear reaction rates was started with the pioneering work of Salpeter \\cite{salp}, who discussed both the extreme cases of \"weak\" and \"strong\" screening, providing suitable expressions for the screening factors \\begin{equation} \\label{eqf}f_{ij}= \\langle\\sigma v\\rangle _{ij,plasma} / \\langle \\sigma v \\rangle _{ij,bare} \\, . \\end{equation} The solar core is  not far from the weak screening case, however it does not satisfy the usual conditions under which the weak screening approximation holds. This is the reason why the problem has been investigated by several authors, see e.g. \\cite{gdgc,koonin,mitler,dzitko,ricci,liolios}.  Gruzinov and Bahcall \\cite{gb98} calculated the electron density in the vicinity of fusing nuclei using the partial differential equation for the density matrix that is derived in quantum statistical mechanics. Their numerical result agrees, within small uncertainties, with Salpeter's weak screening formula. Furthermore, Bahcall et al. \\cite{b2000a} recently provided several arguments that  demonstrate the validity of the Salpeter formula near the solar center with insignificant errors.  The conclusions of Gruzinov and Bahcall \\cite{gb98} are not unanimously accepted. According to Shaviv and Shaviv \\cite{s2}, the weak screening formula does not hold in the Sun.  Some nuclear reactions  are enhanched by the surrounding plasma whereas some others are suppressed. According to  Tsytovitch and Bornatici \\cite{tsy2000,bornatici} a kinetic description of collective plasma effects results in a decrease of the thermonuclear reaction rates in contrast to the Salpeter's enhancement.  We observe that solar models are built by using stellar evolutionary codes which include specific expressions for the the nuclear reaction rates. If one uses  different formulas for the screening factors $f_{ij}$ one obtains different solar models. On the other hand, helioseismology provides precise information on the  sound speed profile and on the properties of the convective envelope, see e.g. \\cite{eliosnoi}, which have to be reproduced by the correct solar model. The main purpose of this paper is to test the screening models by means of helioseismology. We build several solar models corresponding to different screening factors and  compare the results with helioseismic data.  We also comment on the predicted neutrino fluxes. Many of the attempts to modify the screening factors have been produced as an effort to avoid or mitigate the so called solar neutrino puzzle, by reducing the predicted $^8B$ neutrino flux. We will show that a reduction of the screening factors does not  generally imply a reduction of the $^8B$ neutrino flux. ",
        "conclusions": "\\label{conclusions}  We recall here the main points of our discussion:  \\noindent i)The anti-screening predictions of ref \\cite{tsy2000,bornatici} can be excluded by means of helioseismology, since both $u$ and $R_b$ are significantly altered.  \\noindent ii) We find that the  a  no-screening  solar model is not completely consistent with helioseismic data on $u$ and $R_b$, in other words  the existence of a screening effect can be proved by means of helioseismology."
    },
    "0011/astro-ph0011306_arXiv.txt": {
        "abstract": "Selfconsistent magnetospheres of rotating cosmic magnets (neutron stars/pulsars) with arbitrary inclination of the magnetic against the rotation axis are considered. Present studies concentrate on the regime dominated by the force--free surface (FFS). A macroscopic fluid description is applied and radiation reaction is taken into account. As in earlier work of our group, a 'standard set of parameters' is used. Under these conditions, the following features are found among other results: global charge separation exists for all degrees of inclination of the magnetic against the rotation axis; clouds of different charge are seperated by regions of vanishing particle number density; as expected, test particles inserted into the latter regions propagate into one of the adjacent clouds; strong polodial currents exist; locally averaged particle energies for protons typically range up to $10^{16}-10^{17}$ eV, depending on the angle of inclination. ",
        "introduction": "Since the discovery of pulsars in 1968 \\cite{hewish} and their interpretation as a rapidly rotating magnetized neutron stars in the same year \\cite{gold, pacini}, these compact objects are under discussion as powerful accelerators of ultra high energy cosmic ray particles.  From the very beginning of neutron star physics, work on the dynamics of electrically charged particles accelerated in the corresponding electromagnetic fields proceeded on two stages: (1) test particle dynamics in the vacuum fields of rotating magnets, e.g.~\\cite{ostriker, gunn}, demonstrating fundamental mechanisms and (2) self consistent plasma dynamics, e.g.~\\cite{goldreich}, reproducing certain aspects of the structure and evolution of neutron star magnetospheres.  Since then, numerous papers have been published, investigating these matters in great detail, which we will not be able to discuss or even just to mention in this introduction to our present paper. Here we shall concentrate on stage (2), on relativistic plasma dynamics in a regime governed by the force--free surface of a homogeneously magnetized, rapidly rotating sphere with parameters typical for neutron stars.  The notation of a force--free surface (FFS) refers to the dynamic of an electrically charged (test) particle within given electromagnetic fields. By definition the FFS is generated by those points of configuration space at which the Lorentz--force acting on that particle through given electromagnetic fields vanishes. While in published literature a particle that happens to be at such a point often referred to as 'force-free' (ff), we prefer -- in view of the presence of other types of electromagnetic forces (radiation reaction forces) -- to speak of a Lorentz--force--free (Lff) particle in that situation. If ${\\bf B} \\neq 0$ and ${\\bf E} \\neq 0$ \\footnote{Boldface letters are for vectors in 3-dimensional euklidean space. $(\\cdot,\\cdot)$ denotes the scalar product, $[\\cdot,\\cdot]$  the vector product}, as is the case in fields considered here, a particle is  Lff for ${\\bf E} + [\\Bb , {\\bf B}] = 0$, i.e., for (I) $(\\bf{E}, \\bf{B}) = 0$ and (II) $(\\bf{E}, \\Bb)= 0$ and (III) $([\\bf{E},\\bf{H}], \\bm{\\beta}) = |\\bf{E}^2|$. ${\\bf E}$ is the electric field vector, ${\\bf B}$ the vector of the magnetic induction  and $\\Bb$  is the velocity in units of the velocity of light. For $| \\bf{E} | \\ll | \\bf{B} |$, as is the case under premises adopted here, one may expect the second and the third conditions to be inherently fulfilled to some approximation so that the FFS then is caracterized solely by the first condition, $(\\bf{E}, \\bf{B}) = 0$.   From the early works of \\cite{jack1, jack2} magnets rotating in the vacuum with the vector of magnetic dipole moment inclined against their respective rotation axis are known to create such FFS, of which some segments can act as particle traps and thus may have strong bearings on the formation of a neutron star magnetosphere, at least within a certain range of distance from its surface.  In what follow, it will be useful to distinguish the special case of {\\em aligned rotators}, i.e.~rotating magnets with the magnetic axis parallel to the rotation axis ({\\em parallel rotators}) or antiparallel to the rotation axis ({\\em antiparallel rotators}), from the general case of {\\em inclined rotators} and from rotating magnets with the magnetic axis orthogonal to the rotation axis ({\\em orthogonal rotators}). On stage (2), a considerable number of investigations has been published on aligned rotators. Some of these papers will be mentioned below. But only few are devoted to inclined or orthogonal rotators. In the latter case, obviously, considerable formal and numerical complications arise from the lack of rotational symmetry. Analytical approaches were used, for example, by \\cite{jack2, endean, uchida, shibata}.  In the special case of aligned rotators, due to axial symmetry, theoretical results are achieved much easier, even on stage (2). Analytical methods have been applied to the structure of the magnetosphere of aligned rotators for example by \\cite{goldreich, jack1, mestel, fritz1, fritz2, kabu1, kabu2, rylov1, rylov2}. Also, numerical studies on that matter have been performed by \\cite{kuo1, kuo2, krause, ertl, zachariades, neukirch}.   On stage (1) in an earlier work of our group \\cite{wolf89} the authors have integrated numerically the equation of motion for individual test particles within the regime of the FFS and with no restrictions on the relative orientation of the rotational towards the magnetic axis. Results confirm that velocity components orthogonal to the magnetic field vector are efficiently damped by radiation reaction. Thus, test particles tend to follow magnetic field lines, as suggested earlier by \\cite{jack1}. In a certain class of orbits they oscillate about the FFS, while moving along magnetic field lines. The amplitude of these oscillations decreases through radiation losses. Ultimately, in the subsequent regime of lower energy (and on a much larger time scale) particles become subject to drift in azimuthal direction.  On stage (2) in a second paper of our group \\cite{wolf94} the authors have introduced a numerical iterative approach to reproduce sequences of quasi stable plasma configurations forming under the influence of the FFS. In iterative steps charged particles were allowed to be ejected from surface elements of the rotating sphere in quantities locally proportional to the magnitude of the electric vector component normal to the respective surface element. These particles were then allowed to move freely along the appropriate magnetic field lines and to settle down where the projection of the electric onto the magnetic vector vanishes. Thus, {\\em without making use of the equation of motion}, co-rotating, quasi stable, charge separated clouds were reproduced, in consistency with earlier results mentioned above.  In our present work we proceed one step further taking into account in a selfconsistent way virtually {\\em all effects of relativistic particle dynamics}, including radiation reaction and effects of special relativity as, for example, retardation. Our numerial approach is designed  to describe the evolution of locally averaged particle densities, since velocity dispersion is not taken into account. Again, a magnetosphere is allowed to build up from the initial vacuum through particle ejection from the spherical surface, similar to the procedure described above. Here particles are allowed to be ejected from surface elements, given an appropriate direction of the electric vector,in quantities locally proportional to the magnitude of the electric vector component projected onto the magnetic field line.  Thereby, we intend to clearify, on stage (2), the evolution, selfconsistent structure and stability properties of plasma configurations forming within the regime of the FFS of a rotating cosmic magnet with no restrictions on the relative orientation of the rotational towards the magnetic axis. Also, we want to evaluate mean energy values locally achieved by particles in that regime.  As in earlier works of our group, e.g.~\\cite{laueth, th91, th94}, we  apply a model represented by a rotating, ideally conducting, homogeneously magnetized sphere with arbitrary inclination of the magnetic against the rotation axis. A 'standard set of parameters' representing well-known properties of typical neutron stars is attributed to this model: the stellar mass which is taken equal to the solar mass $m_{N} =  m_{{\\rm sun}}$, the stellar radius $r_{N} = 10^6$cm, the angular velocity $\\omega = 20\\,\\pi\\,{\\rm s}^{-1}$, and the magnetic dipole moment $\\mu = 10^{30}$ G ${\\rm cm}^3$ \\footnote{Gaussian units are used throughout this paper. Thus, electric and magnetic field strengths are measured in units of $1 {\\rm G} = 300 {\\rm V/cm}$. For the standard set of parameters, the magnetic field strength is about $B_{p} \\simeq 2 \\cdot 10^{12}$ \\,G and the electricic field strength is approximately $E_p \\simeq 10^{10} G$ in the polar region. Under given parameter values, the radius of the light--cylinder (often referred to as 'light radius') $r_L = \\o/c$, outside of which corotation cannot exist, is $r_L = 5 000$ km, corresponding to almost the radius of the earth, for comparison.}.  In order to investigate the regime of the FFS with an appropriate resolution we concentrate on the near zone of the neutron star up to $20\\,r_N$.  From preceeding estimates as well as from subsequent simulations gravitational forces exerted by the rotating neutron star and by the magnetosphere itself onto individual particles, as well as effects of general relativity were found to be negligible for the standard set of parameters. Contributions to the  Lorentz--force originating from magnetic fields created by magnetospheric particle currents can also be neglected, in agreement with earlier conclusions of \\cite{ertl}. Spontaneous pair creation still turns out to be insignificant, even within the very strong electromagnetic fields of polar regions.   In chapter 2 of what follows, we rediscuss  force--free surfaces associated with rotating magnets. Thereafter, in chapter 3, we display equations of {\\em individual} as well as of {\\em collective particle motion} in terms of a non-neutral two-component fluid description, and we then proceed to a description of appropriate tools for numerical treatment in chapter 4. Results will be given in chapter 5 and subsequently discussed in chapter 6. ",
        "conclusions": "The resulting quasi-stationary magnetospheres are not global force-free. In regions with vanishing particle densities the vacuum electromagnetic field is approximately undisturbed by the non-neutral plasma. In regions with high particle number densities the  projection of the electric vector onto the tangent to the magnetic field vanishes nearly, but not complete. Small deviations from a force--free situation leads due to the extrem strong electromagnetic fields immediately to high particle energies. All in all, independent of the inclination angle, highly relativistic plasmas are found, which leads to the necessity to take the radiation reaction into account.  In the case of the parallel and antiparallel rotator, due to the fact that we did not found particles with low oder moderate particle energies, the usage of the ultra--relativistic approximation of the Lorentz-Dirac-Landau-equation by \\cite{ertl} and \\cite{zachariades} in their numerical calcualtion is justifiable, shown by our studies. Futhermore, we can confirm the existence of clouds with different charges, seperated by regions of vanishing particle number density (vacuum--gaps), also found by \\cite{krause}, \\cite{ertl}, \\cite{zachariades}, \\cite{wolf94} and \\cite{neukirch}. All these works including ours found no global force-free magnetosphere. Force-free magnetospheres are often used as an assumption in analytical works. The structure of the quasi-stationary magnetospheres predicted by \\cite{jack1} with analytical models can be confirmed with our studies, especially the importance of the FFS. Nevertheless, the closed polodial currents proposed by \\cite{jack1} can not be proven by our studies due to the limited simulation volumen of 20 $r_N$. The density distribution of the plasma is described very well by only one quadrupole mode, confirming with \\cite{goldreich} in their amplitude and power law.  In the case of the inclined and orthogonal rotator only few literature is published. The by \\cite{ endean} proposed approach solving the stationary Valsov-equation and the Maxwell-equations selfconsistently led to an inside the light-cylinder corotating, charge seperated magnetosphere. The charge separation independent of the inclination angle was also found by \\cite{wolf94}, which can be confirmed by our present studies. Basicly, the structure of the quasi-stationary magnetospheres in our calculations confirms with those found by \\cite{jack2} analytically. Furthermore, we can attest the dependency of the electric monopole of the rotating sphere of the declination angle, first proposed by \\cite{jack2}. Beside this, for the first time with our work statements on particle engeries, currents and drifts based on numerical, full dynamical studies are possible. Due to the limited simulation volume we can not prove if outside the light-cylinder closed currents are formed (starting and ending on the surface of the sphere). The knowledge of the global current system is important due to the fact that the $[{\\bf E, B}]$-drift is important at long time scales. As a consequence of these force charged particles move back to the surface and may change the electric monopol of the sphere. The structure of the FFS and resulting from this the structure of the magenetosphere is depending on the electric monopol of the sphere.  The high particle energies in the existing polodial currents are relevant for induced pair production. In the present paper we were able to ignore the resulting $e^{\\pm}$--plasma. Nevertheless, relativistic currents in these plasmas may cause microscopic instabilities which could explain the non--thermal radiation in real pulsar  magnetospheres. Analytical studies on this topic (e.g.~\\cite{lyutikov}) assume Lor\\-entz--factors in these currents of $\\gamma \\approx 10^7$ and particle densities in the scale of the  Goldreich \\& Julian--density, as were confirmed by our studies.  Regarding the discussion about neutron stars as cosmic accelerators for ultra high energy cosmic ray particles the very high particle energies proven by our studies are remarkably. I.e., in the case of the $120^{\\circ}$--rotator we found averaged proton energies up to $10^{17}$ eV. In the case of higher magnetic fields, than given by the 'standard set of parameters' higher particle energies are possible. But, investigating these cases one schould prove whether a classical approach is applicable. It is important to note, that due to the limited simulation volume used in our work we can not predict if these high energy particle are able to leave the neutron star magnetoshphere and if, at which particle energies."
    },
    "0011/astro-ph0011312_arXiv.txt": {
        "abstract": "We show that if the dwarf-nova disc instability model includes the effects of heating by stream impact and tidal torque dissipation in the outer disc, the calculated properties of dwarf-nova outbursts change considerably, and several notorious deficiencies of this model are repaired. In particular: (1) outside-in outbursts occur for mass transfer rates lower than in the standard model as required by observations; (2) the presence of long (wide) and short (narrow) outbursts with similar peak luminosities is a natural property of the model. Mass-transfer fluctuations by factors $\\sim$ 2 can explain the occurrence of both long and short outbursts above the cataclysmic variable period gap, whereas below 2 hr only short normal outbursts are expected (in addition to superoutbursts which are not dealt with in this article). With additional heating by the stream and tidal torques, such fluctuations can also explain the occurrence of both outside-in and inside-out outbursts in \\object{SS Cyg} and similar systems. The occurrence of outside-in outbursts in short orbital-period, low mass-transfer-rate systems requires the disc to be much smaller than the tidal-truncation radius. In this case the recurrence time of both inside-out and outside-in outbursts have a similar dependence on the mass-transfer rate $\\dot{M}_{2}$. \\keywords accretion, accretion discs -- instabilities -- (Stars:) novae, cataclysmic variables -- (stars:) binaries : close ",
        "introduction": "Dwarf novae (DN) are erupting cataclysmic variable stars. The 4 - 6 mag outbursts last from few days to more than a month and the recurrence times can be as short as a few days and as long as 30 years (see e.g. \\cite{w95}). In cataclysmic variables, a Roche-lobe filling low-mass secondary star is losing mass which is accreted by a white dwarf primary. In DNs the matter transfered from the secondary forms an accretion disc around the white dwarf; it is this disc which is the site of the outbursts. It is believed that an accretion disc instability due to partial hydrogen ionization is triggering the outbursts. The disc instability model (DIM; see Cannizzo 1993, Lasota 2000 for reviews) is supposed to describe the whole DN outburst cycle. This model assumes the $\\alpha$-prescription (\\cite{shak73}) for the viscous heating and the angular momentum transport in the disc. In such a framework, the outbursts are due to propagating heating and cooling fronts, while during quiescence the disc is refilling the mass lost during the outburst. In the standard version of the model, the mass transfer rate is assumed to be constant during the whole outburst cycle. In order to reproduce outburst amplitudes and durations, the $\\alpha$-parameter has to be larger in outburst than in quiescence (Smak 1984).  The DIM is widely considered to be the model of DN outbursts because it identifies the physical mechanism giving rise to outbursts, and with the above mentioned assumptions it produces outburst cycles which roughly look like the real thing. A closer look, however, shows that the model fails to reproduce some of the fundamental properties of dwarf-nova outbursts. The main deficiencies of the DIM were recently summarized by Smak (2000). First, models predict a significant increase of the disc luminosity during quiescence (see e.g. Fig. \\ref{fig3}), an effect which is not observed. Second, in many systems the distribution of burst widths is bimodal (van Paradijs 1983). The DIM can reproduce the width bimodality but, contrary to observations, its narrow outbursts have lower amplitudes than the wide ones. Third, observations clearly show that during outbursts the secondary is irradiated and the mass transfer rate increases. The standard DIM does not take this into account. To these problems pointed out by Smak one can add several other deficiencies. The model predicts outbursts beginning in the outer disc regions (`outside-in') only for high mass transfer while observations show the presence of such outbursts at apparently low mass transfer rates. In most cases the model predicts quiescent accretion rate which are at least two orders of magnitude lower than the rates deduced from quiescent luminosities. Finally, it is clear that superoutbursts observed in SU UMa stars and other systems (Smak 2000) cannot be explained by the standard DIM.  All this leads to the obvious conclusion that the standard version of the DIM has to be enriched by inclusion of physical processes which might be important but which, for various reasons, were not taken into account in the original model. Osaki (1989) proposed that superoutbursts are due to a tidal-thermal instability. Smak (2000) suggested that problems of this model could be solved by a hybrid which would combine it with the radiation-induced enhanced mass transfer model (Osaki 1985; Hameury, Lasota \\& Hur\\'e 1997). He included mass transfer enhancements during outburst into the standard model and suggested that this could explain properties of the narrow and wide outbursts which alternate in many systems (Smak 1999a). Hameury, Lasota \\& Warner (2000), extending previous work by Hameury et al. (1997) and Hameury, Lasota \\& Dubus (1999), showed that inclusion of effects such as mass transfer variations, formation of holes in the inner disc and irradiation of both the disc and the secondary allows to reproduce many properties of the observed outbursts.  Some questions, however, were left without satisfactory answers. First, the problem of `outside-in' outbursts occurring for too high mass-transfer rates was still unsolved. Second, Smak (1999a) obtained narrow and wide outbursts by including irradiation induced mass transfer fluctuations but wide outbursts in his model result from fluctuations whose amplitude bring the disc to a stable state. In such a case the width of wide outbursts would not be determined by the disc structure alone. There is nothing fundamentally wrong with such a model and long standstills of Z Cam systems are most probably due to such fluctuations, as proposed by Meyer \\& Meyer-Hofmeister (1983) and shown by Buat-M\\'enard et al. (2000). However, one should check that the model is complete, that no other mechanism can explain outburst width distribution.  In the present paper we study two mechanisms which dissipate energy in the outer parts of an accretion disc. They are: (1) the impact on the outer boundary of the accretion disc of the gas stream which leaves the secondary through the $L_1$ critical point of the Roche lobe; this impact zone is observed in the light-curves as a bright (`hot') spot; (2) the viscous dissipation induced by the tidal torques which are responsible for the disc outer truncation. These mechanisms which may heat up the disc's outer regions have been known for a long time (\\cite{pacz77,pap-pring77,meyer83}), but have been neglected in most of DIM calculations. The exception is Ichikawa \\& Osaki (1992) who included both effects in their model. They used, however, a viscosity prescription specially designed to suppress inside-out outbursts and did not perform a systematic study of the effects as a function of the mass-transfer rate, size of the disc etc.  We add stream-impact and tidal torque heating effects to the Hameury et al. (1998) version of the DIM, presented in Sect. 2. We use a 1D approximation because the dynamical time-scale is shorter than the thermal time-scale. We study the effect of the stream impact (Sect. 3) and the tidal torques (Sect. 4) on outburst properties in short and long period DN systems. Both effects change the energy equation of the standard DIM. In both cases we find similar results: when compared to the standard case, `outside-in' outbursts occur for lower mass transfer rate. The recurrence time is decreased by both the stream impact and dissipation of the tidal torques. We conclude, therefore, that heating of the outer disc by the stream impact and the rate of work by tidal torques provides a satisfactory explanation of the presence of `outside-in' outbursts at realistic mass-transfer rates.  The most important result of our investigation is that when these effects are included one can obtain two types of outbursts: short (`narrow') and long (`wide') (Sect. 5). We find a critical mass transfer rate for which these two type of outburst `alternate'. This mass transfer rate is very close to that deduced from evolutionary calculations for systems with orbital periods longer than 3 hours (`above the period gap'). For lower mass-transfer rates only narrow outbursts are present; for higher ones, outbursts are only of the wide type. Fluctuations of the mass-transfer rate allow to expand the range of orbital periods for which the bimodal outburst width distribution is occurring. We conclude that the observed bimodal outburst width distribution is caused by the outer-disc heating.  In Sect. 6 we discuss conditions for which outside-in outbursts can occur. We show that one can naturally account for the occurrence of both types of outbursts in \\object{SS Cyg} if the mass transfer rate fluctuates by factors $\\lta 2$ around a mean value which is close to what is generally assumed as the average mass transfer rate in this system. Outside-in outbursts in systems below the period gap, which have low mass transfer rates, can be explained only if the disc is significantly smaller than the tidal truncation radius. We show that this is in agreement with observations and consistent with expectations of the tidal-thermal instability model, at least during the first half of the supercycle of SU UMa systems. We also point out the ambiguity in observational determination of outburst `outside-in' or `inside-out' type, also respectively called A and B types. These types are generally assigned to outbursts according to their shape and the UV delay. We question the general validity of such criteria and argue that the eclipse profile observations during outbursts is the only reliable diagnostics of the outburst type. ",
        "conclusions": "We have shown that the inclusion of additional heating effects such as the stream impact and tidal dissipation might solve several problems of the standard DIM.  The most important result is that when the tidal torque dissipation in the outer disc regions is included, one obtains for a certain range of mass-transfer rates an alternation of short and long outbursts which is observed in long orbital period system such as \\object{SS Cyg} (\\cite{can92}) or \\object{Z Cam} (\\cite{oppen98}). The long outbursts are slightly brighter than short outbursts as observed. Moderate fluctuations of the mass-transfer rate would extend the range for which such alternating outbursts are present. For systems with mass transfer rates close to the stability limit, such fluctuations will produce standstills as shown in Buat-M\\'enard et al. (2000). This solution of the long/short outburst problem (Smak 2000) is different from the one suggested by Smak (1999a). In our model mass-transfer fluctuations play only an auxiliary role and long outbursts do not result from fluctuations bringing the disc to a steady state.  We showed that additional heating effects make the outside-in outburst possible at low mass transfer rates without changing the $\\alpha$-prescription. Our results on mass transfer rates and outburst types are in good agreement with most estimates in the case of long period systems. In the case of short period systems of the SU UMa type, for outside-in outburst to occur the disc would have to be smaller than usually assumed. For small discs the recurrence times of both the inside-out and outside-in outbursts is $\\dot{M}_2$-dependent. It must be stressed out that there are still many uncertainties on the observational determination of outburst types. For instance, the UV delay is not necessarily a good criterion for determining the outburst type; the eclipse profile during outburst is a much better criterion, but unfortunately restricted to a small subset of dwarf novae.  There are still several points that need to be clarified. First, the stream impact effect geometry is not well known and is considered here in a 1D approximation (as for the tidal dissipation). We do not know how far and how efficiently the energy propagates in the disc. The expression of the tidal torque can also be subject to discussion. Moreover, we did not include irradiation here, and it would be interesting to study the combined effects of all phenomena. Finally, as raised above, the existing discrimination criteria on the outburst types A and B have to be tested and improved."
    },
    "0011/astro-ph0011254_arXiv.txt": {
        "abstract": "We report the first wide-field mapping of the kinematics and stellar populations in the E3 galaxy NGC\\,4365. The velocity maps extend previous long-slit work. They show two independent kinematic subsystems: the central $300 \\times 700$~pc rotates about the projected minor axis, and the main body of the galaxy, $3\\times4$~kpc, rotates almost at right angles to this. The line-strength maps show that the metallicity of the stellar population decreases from a central value greater than solar, to one-half solar at a radius of 2~kpc. The decoupled core and main body of the galaxy have the same luminosity-weighted age, of $\\approx$14~Gyr, and the same elevated magnesium-to-iron ratio. The two kinematically distinct components have thus shared a common star formation history. We infer that the galaxy underwent a sequence of mergers associated with dissipative star formation that ended $\\gta$12~Gyr ago. The misalignment between the photometric and kinematic axes of the main body is unambiguous evidence of triaxiality. The similarity of the stellar populations in the two components suggests that the observed kinematic structure has not changed substantially in 12~Gyr. ",
        "introduction": "\\label{sec:intro}  The existence of decoupled cores in $\\approx$30\\% of the early-type galaxies is strong evidence that mergers play an important part in the evolution of these systems \\citep[\\eg\\/][]{zf1991}. Most likely, decoupled cores originate from the accretion of material with angular momentum misaligned from that of the main galaxy. A few galaxies have been studied in sufficient detail to explore when that material was accreted or whether the event was associated with gaseous dissipation and star formation. For example in the morphologically normal galaxies IC\\,1459 \\citep{fi1988} and NGC\\,5322 \\citep{ben1988,rw1992} there is no difference between the colors or line-strengths of the stellar populations in the decoupled core and the main galaxy. In contrast, the disturbed shell galaxy NGC\\,2865 shows evidence for very recent star formation in the decoupled component \\citep*{hcb1999}.\\looseness=-2  NGC\\,4365 is one of the first elliptical galaxies in which minor axis rotation was discovered \\citep*{wbm1988,bsg1994}.  \\citet[][hereafter SB]{sb1995} deduced the remarkable kinematic structure of NGC\\,4365 from three long-slit spectra. The main body of the galaxy rotates around its {\\em major}\\/ axis, reaching a maximum velocity of $\\approx$50~\\kms, whereas at smaller radii ($r=$ 2--3\\arcsec) the peak rotation velocity is 80~\\kms\\/ around the {\\em minor}\\/ axis. SB found that the decoupled core is flatter than the main galaxy (ellipticity $\\varepsilon_{\\rm core}$=0.39 cf.\\ $\\varepsilon_{\\rm main}$=0.23). By carrying out a careful double Gaussian analysis of the kinematics, SB estimate that $V/\\sigma= 1.3$ for the core, consistent with a disk or bar. They deduced that the rotating core contributes 25\\% of the light within $r\\approx6$\\arcsec, and accounts for roughly 2\\% of the {\\em total}\\/ mass. The giant elliptical galaxy in which it is embedded has a central velocity dispersion of 275~\\kms. They found that the decoupled core contains a high-metallicity population, with an enhanced [Mg/Fe] ratio as commonly found in giant ellipticals. SB concluded that NGC\\,4365 is triaxial, and that the formation of the core involved substantial gaseous dissipation and star formation. However, they could not distinguish formation scenarios where the merger(s) occur early ($z> 2$-3) from those where the mergers occur more recently ($z<1$).\\looseness=-2  \\begin{figure*} \\epsfig{file=col_fig.eps} \\caption[]{\\label{fig:color}Maps of a) surface brightness as reconstructed from our data, b) mean streaming velocity $V$, c) velocity dispersion $\\sigma$, d) \\mgb\\/ line-strength, and e) H$\\beta$ line-strength, for NGC\\,4365. The maps are based on two partially overlapping \\sauron\\/ pointings of 4$\\times$1800~s each, sampled with $0\\farcs8\\times0\\farcs8$ pixels. Isophotal contours of $\\mu_V = 18$ and 19 mag/arcsec$^2$ are overplotted in panels (a) to (e) (dashed lines). A zero-velocity contour line (solid black line) is shown in panel (b). The enlarged core region of the velocity map indicates the regions which we have used for our line-strength analysis: The central point is indicated by the square, the decoupled core by the sectors along the major axis and the main body at the same radius by the sectors along the minor axis.} \\end{figure*}  HST images reveal that NGC\\,4365 has a smooth intensity profile with no signs of dust, an average ellipticity of 0.26, and a very shallow central cusp. The isophotes are disky for $1\\arcsec<r<4$\\arcsec\\/ and boxy at larger radii. At the very centre ($r \\lta 0\\farcs2$) there is a blue point source \\citep{df1995,car1997}. Deep ground-based images show no evidence for shells or other morphological peculiarities out to $\\approx 4 \\, r_e$ (Blakeslee, private communication).  Surface-brightness-fluctuation measurements indicate that NGC\\,4365 is in the Virgo W cloud beyond the main cluster \\citep*{jtl1998}. We take the distance modulus to be 31.7 mag, a distance of 22~Mpc, so that 1\\arcsec\\/ is $\\approx$100~pc. The effective radius $r_e$ is 57\\arcsec, or $\\approx$5.7~kpc \\citep{bur1987}.  In this Letter, we present the first complete maps of the kinematics and stellar populations of NGC\\,4365, taken with the wide-field integral-field spectrograph \\sauron\\ \\citep{bac2001}. In \\S2, we briefly describe the observations. The kinematics and line-strength index maps are presented in \\S3, where we consider age-metallicity diagnostics, the Mg-$\\sigma$ diagram, and the spatial distribution of non-solar abundance ratios. We discuss the implications of these results for formation scenarios in \\S4, and summarize our conclusions in \\S5. ",
        "conclusions": "The \\sauron\\ maps present a complete view of the kinematics and stellar populations of NGC\\,4365. They show two independent kinematic subsystems: the central $300 \\times 700$~pc and the main body of the galaxy, rotating almost at right angles to each other. The misalignment of $82^{\\circ}$ between the photometric and kinematic axes of the main body is unambigous evidence of triaxiality. The \\sauron\\ maps enable us to compare the stellar population in the decoupled component with that in the main body of the galaxy at the same radius. We find these populations to be indistinguishable in age, metallicity and abundance ratios. We find an age of $\\approx$14~Gyr and a decrease in metallicity, from larger than solar in the center to half solar at a radius of 2~kpc ($\\approx0.4 \\, r{_e}$). We suggest that NGC\\,4365 underwent dissipative star formation at high redshift, probably through one or more mergers.  Later generations of stars formed a more centrally-concentrated, metal-enriched stellar population. Star formation was complete and the residual gas was exhausted roughly 12~Gyr ago. This also suggests that the observed kinematics and triaxial structure is stable."
    },
    "0011/astro-ph0011548_arXiv.txt": {
        "abstract": "Dynamical studies of MHD turbulence on the one hand, and arguments based upon magnetic helicity on the other, have yielded seemingly contradictory estimates for the $\\alpha$ parameter in turbulent dynamo theory. Here we show, with direct numerical simulation of three-dimensional magnetohydrodynamic turbulence with a mean magnetic field, $\\OB$, that the constraint on the dynamo $\\alpha$-effect set by the magnetic helicity is time-dependent. A time-scale $t_c$ is introduced such that for $t<t_c$, the $\\alpha$-coefficient calculated from the simulation is close to the result of Pouquet \\etal and Field \\etal, $-\\frac{\\tau_{cor}}{3}( \\langle \\v \\cdot \\nabla \\times \\v \\rangle  - \\langle \\b \\cdot \\nabla \\times \\b \\rangle )$; for $t>t_c$, the classical result of the $\\alpha$-coefficient given by the Mean-Field Electrodynamics is reduced by a factor of $1/\\left( {R_m |\\OB|^2/v_{rms}^2} \\right)$, as argued by Gruzinov \\& Diamond, Seehafer and Cattaneo \\& Hughes. Here, $R_m$ is the magnetic Reynolds number, $v_{rms}$ the rms velocity of the turbulence, $\\tau_{cor}$ the correlation time of the turbulence, and $\\overline B$ is in velocity unit. The applicability of and connection between different models of dynamo theory are also discussed. ",
        "introduction": "The generation and amplification of magnetic field in many astrophysical systems are often attributed to the turbulent dynamo effect. Since the seminal papers by Parker (1955) and Steenbeck, Krause \\& R\\\"adler (1966), a whole body of theory, namely, {\\it Mean Field Electrodynamics} (MFE) has been developed to explain the dynamics of magnetic field generation by helical turbulence in a conducting fluid (Moffatt 1978, or Krause \\& R\\\"adler 1980). In such a fluid, the velocity field, $\\v$, stretches the magnetic field, $\\b$, in such a way that the correlation of $\\v$ and $\\b$ results in an electromotive force, $ \\langle \\v \\times \\b \\rangle $ that amplifies the mean (large-scale) magnetic field, $\\OB$, through the relation \\be \\langle \\v \\times \\b \\rangle  = \\alpha \\OB. \\ee Here, the coefficient $\\alpha$ represents the so-called $\\alpha$-effect (Moffatt 1978); it is calculated in MFE as $\\alpha_{MFE} \\sim -\\frac{\\tau_{cor}}{3} \\langle \\v \\cdot \\nabla \\times \\v \\rangle $, where $\\tau_{cor}$ is the correlation time of the turbulence. MFE is a kinematic theory, in that the velocity field is prescribed and no back reaction of the magnetic field on the velocity field is considered. Therefore, its applicability to circumstances where the velocity field is affected by the growing magnetic field is questionable. Several authors have extended MFE to include the quenching of the $\\alpha$-effect due to the back reaction of the magnetic field. By numerically solving the spectral MHD equations using a closure method known as the EDQNM (Eddy-Damped Quasi-Normal Markovian) approximation, Pouquet, Frisch \\& L\\'eorat (1976, hereafter PFL) find that $\\alpha=\\int_k \\alpha_k dk$, where $\\alpha_k$ is determined by, in Fourier space, the difference between the spectra of the kinetic helicity correlation function, $\\langle \\v \\cdot \\nabla \\times \\v \\rangle$, and the current helicity correlation function, $\\langle \\b \\cdot \\nabla \\times \\b \\rangle$. A similar result was found by Field, Blackman \\& Chou (1999, hereafter FBC), who consider the back reaction of magnetic field by treating $\\v$ and $\\b$ on an equal footing, and give a model in which the $\\alpha$-coefficient (for later purposes, we call it $\\alpha_1$) can be expressed as \\be \\alpha_1 = -\\frac{\\tau_{cor}}{3}\\left( \\langle \\v \\cdot \\nabla \\times \\v \\rangle - \\langle \\b \\cdot \\nabla \\times \\b \\rangle \\right) \\ee for relatively small values of $\\left({\\overline B}/{v_{rms}}\\right)^2$.  However, the nonlinear nature of the problem introduces so much difficulty that the effect of the back reaction is still under debate. One of the objections to the application of $\\alpha_1$ to the galactic dynamo has its root in the problem of large magnetic Reynolds number and small-scale fields. Vainshtein \\& Cattaneo (1992, see also Cattaneo \\& Vainshtein 1991) argue that for systems with large $R_m$, $\\alpha$ is reduced by a factor of $R_m$ from its kinematic value, i.e., \\be \\alpha \\propto \\frac{1}{R_m}, \\ee as the small-scale magnetic field grows quickly to turn off the generation of magnetic flux at large scales. These authors applied the conservation of the square vector potential of 2D ideal MHD. The conservation of magnetic helicity of 3D MHD in steady state was later studied by Seehafer (1994, 1995) who provided another model of the dynamo $\\alpha$-effect in which the $\\alpha$-coefficient (for later purposes, we call it $\\alpha_2$) is given by \\be \\alpha_2 = -\\frac{\\la \\langle \\b \\cdot \\nabla \\times \\b \\rangle}{{\\overline B}^2}. \\ee Because $R_m = Lv_{rms}/\\la$, $\\alpha_2 \\propto 1/R_m$ also. Because $R_m \\gg 1$ in astrophysical systems such as the Galaxy, both relation (3) and relation (4) suggest strong suppression of the dynamo $\\alpha$-effect.  Another model of the dynamo $\\alpha$-effect was introduced by Gruzinov \\& Diamond (1994, 1995, 1996) by studying how the conservation of magnetic helicity affects the result like relation (2). They realized that because Ohmic dissipation of the current helicity $\\langle \\b \\cdot \\nabla \\times \\b \\rangle$ in (2) changes the magnetic helicity, the dynamics of the latter will affect $\\alpha_1$, which should be modified to (for later purposes, we call it $\\alpha_3$) \\be \\alpha_3 =  \\frac{\\alpha_{MFE}}{1+R_m ({\\overline B}/v_{rms})^2}. \\ee Since $R_m$ is usually very large in astrophysics, this again implies that $\\alpha$, in steady state, is far smaller than its classical value, so that it is too small to be important for the generation of large scale magnetic fields. Note that in the limit of $R_m \\gg 1$ and $v_{rms} \\sim b_{rms}$, $\\alpha_3$ is reduced to $\\alpha_2$.  $\\alpha_2$ and $\\alpha_3$ are different from $\\alpha_1$ in the following two aspects: first, $\\alpha_1$ does not explicitly show strong suppression of the $\\alpha$-effect for large $R_m$; second, neither of the dynamical studies of PFL and FBC, which give $\\alpha_1$, explicitly considered the conservation of magnetic helicity that led to $\\alpha_2$ and $\\alpha_3$. Under what conditions the magnetic helicity constraint, which is essential in the models of $\\alpha_2$ and $\\alpha_3$, enters the derivation of $\\alpha_1$ is still an open question. In other such words, if such magnetic helicity constraint can be relaxed in systems of large $R_m$, which of these three models remains valid? To provide insight into this problem, in the following we study the dynamics of magnetic helicity using the magnetic helicity conservation equations, and the time dependence of $\\alpha$ using a numerical simulation of MHD turbulence. We find that both the magnetic helicity development and the $\\alpha$-coefficient are time dependent. We find that the classical MFE result is valid up to a critical time that we calculate, and the magnetic helicity constrained result is valid thereafter. Which value to apply therefore depends on the circumstances.  This paper is presented in the following structure: in Section 2 we provide a model for the time-dependence of the magnetic helicity development, and introduce a critical time, $t_c$, to separate the two important stages of the magnetic helicity evolution; in Section 3.1 we present our numerical model that is used to study the time-dependence of the dynamo $\\alpha$-effect; the numerical results are given in Section 3.2; in Section 4 we discuss the implications of our analytic model and the numerical results, which are applied to the Galactic dynamo; conclusions are given in Section 5. ",
        "conclusions": "We studied the constraint of magnetic helicity on the dynamo $\\alpha$-effect with 3D direct numerical simulation under periodic boundary conditions. The dynamics of magnetic helicity affects the dynamo $\\alpha$-effect only after the magnetic helicity at small scales is built up and the magnetic helicity dynamics enters a stationary state. Such building-up process can be understood as a pumping effect of the dynamo $\\alpha$-effect, and the $\\alpha$-coefficient during this non-stationary pumping stage can be estimated according to the model by Pouquet \\etal (1976) or Field \\etal (1999). As the small-scale magnetic helicity is built up to the level limited by the realizability condition, the $\\alpha$-effect is quenched, as suggested by Gruzinov \\& Diamond (1994, 1995, 1996) and Cattaneo \\& Hughes (1996). The $\\alpha$-coefficient during such magnetic helicity constraining stage can be estimated according to the model by Seehafer (1994, 1995, see also Blackman \\& Field 2000). A critical time, $t_c \\sim t_{eddy}\\left(b_{rms}/{\\overline B}\\right)^2$, is introduced to separate these two stages."
    },
    "0011/astro-ph0011062_arXiv.txt": {
        "abstract": "CorMASS (Cornell Massachusetts Slit Spectrograph) is a compact, low-resolution ($R\\sim300$), double-pass prism cross-dispersed near-infrared (NIR) spectrograph in operation on the Palomar Observatory 60-inch telescope.  Its 2-dimensional spectral format provides simultaneous coverage from $\\lambda\\sim 0.75 \\micron$ to $\\lambda \\sim 2.5 \\micron$ (\\textit{\\'{z}JHK} bands).  A remotely operated cold flip mirror permits its NICMOS\\,3 detector to function as a $K_s$ slit viewer to assist object placement into the $2\\arcsec \\times 15\\arcsec$ slit.  CorMASS was primarily designed for the rapid spectral classification of low-mass stellar and sub-stellar objects identified by the Two-Micron All Sky Survey (2MASS).  CorMASS' efficiency and resolution also make it a versatile instrument for the spectral observation and classification of many other types of bright objects ($K<14$) including quasars, novae, and emission line objects. ",
        "introduction": "Recent successes by deep near-infrared (NIR) surveys such as the Two Micron All Sky Survey (2MASS; \\citep{skr97}), the Deep Near-Infrared Survey (DENIS; \\citep{epc97}), and the Sloan Digital Sky Survey (SDSS; \\citep{gun95}) in identifying very low-mass stellar and sub-stellar field candidates have made a tremendous impact upon the study of low-mass objects in the local stellar neighborhood.  Spectroscopic follow-up of low-mass survey candidates has revealed objects with atmospheric features different from the TiO and VO features seen in the atmospheres of M-dwarfs.  Two new classifications, L with $1300 \\la T_{eff} \\la 2000$ \\citep{ki00a} and T with $T_{eff} \\la 1200-1300$ \\citep{feg96,bur99,ki00a}, have been proposed to account for the unique spectral characteristics of these objects at and below the bottom of the main sequence \\citep{mar97,kir99}.  L-dwarf spectra are dominated by metallic hydrides and alkali atomic features as the oxides of M-dwarfs condense into solids at these cooler atmospheric temperatures.  T-dwarfs show the distinctive ${\\rm CH}_4$ absorption features at $1.6 \\micron$ and $2.2 \\micron$.  Gl 229B is the archetypical T-dwarf \\citep{nak95,opp98}.  Motivated by the potential identification of a few $10^3$ candidate L and T-dwarfs in the northern sky by 2MASS \\citep{rei99}, we have built a spectrograph specifically for the efficient confirmation and classification of low-mass object candidates.  This paper describes the Cornell Massachusetts Slit Spectrograph (CorMASS) and presents a sample of spectral observations from its first runs on the Palomar 60-inch telescope.  CorMASS is a double-pass prism cross-dispersed NIR grating spectrograph. Simultaneous wavelength coverage spans from $\\lambda\\sim 0.75 \\micron$ to the $256 \\times 256$ HgCdTe NICMOS\\,3 array cut-off at $\\lambda\\sim 2.5 \\micron$ (the \\textit{\\'{z}JHK} photometric bands)\\footnote{\\textit{\\'{z}} is a SDSS photometric band with $\\lambda_c = 0.91\\micron$ and $\\Delta\\lambda = 0.12\\micron$.  The long wave cut-off is determined by CCD sensitivity \\citep{fuk96}.}.  This wavelength coverage is dispersed into 6 orders on the array.  A spectral resolution of $R\\equiv\\lambda/\\Delta\\lambda\\sim 300$ maximizes signal from relatively faint objects, while keeping resolution high enough to identify features of the target between blended airglow lines, including in H-band.  An externally driven cold flip mirror permits the reflective slit to be imaged in $K_s$ on the NICMOS\\,3 array using a complementary optical train.\\footnote{$K_s$ is a passband created for the 2MASS survey by one of us (M. Skrutskie). This band is shortened and shifted blue-ward compared to the conventional \\textit{K}-band to reduce in-band thermal emission \\citep{per98}.} The slit viewer's $35\\arcsec$ field-of-view (FOV) enables rapid object identification and accurate placement of science targets into the slit for spectroscopy. The compact instrument optical and mechanical design occupies an ${\\rm LN}_2$ cooled volume of $14.5\\arcsec \\times 6.5\\arcsec \\times 6\\arcsec$ inside the Dewar.  After presenting our scientific requirements for the instrument in \\S 2, the optical design is discussed in \\S 3. The mechanical design is discussed in \\S 4. Electronics and Data Acquisition are covered in \\S 5.  In \\S 6 we discuss performance, and in \\S 7 data reduction is discussed and sample spectra are presented. ",
        "conclusions": "This paper described the newly built and commissioned NIR low-resolution ($R\\sim300$) prism cross-dispersed spectrograph CorMASS in operation on the Palomar 60-inch telescope. This instrument was designed primarily for the identification of late-M and L-dwarf spectral features.  CorMASS spectra of low-mass objects from the L, M, and T-dwarf classes, as well as the planetary nebula NGC 7027, were presented. Instrument characteristics include:\\\\ \\indent1. Low-cost compact design.\\\\ \\indent2. Simultaneous $\\lambda$ coverage of the \\textit{\\'{z}JHK} NIR bands with R $\\sim 300$.\\\\ \\indent3. A two-position flip mirror for imaging an alternate slit viewing train.\\\\ \\indent4. Spectrograph sensitivity of $K=14.0$ in $t_{int}=3600$ sec for S/N=5. \\\\"
    },
    "0011/astro-ph0011581_arXiv.txt": {
        "abstract": "We report VLA \\HI\\/ Zeeman observations toward the W49A star-forming region and the SNR W49B.  Line of sight magnetic fields (\\Blos\\/) of 60 to 300 \\muG\\/ at \\tf\\/ resolution were detected toward W49A at velocities of \\4\\/ and \\7\\/.  The \\Blos\\/ values measured toward W49A show a significant {\\em increase in field strength with higher resolution} especially for the \\4\\/ \\HI\\/ component.  The \\HI\\/ gas in the velocity range $-5$ to 25 \\kms\\/ toward W49A shows good agreement both kinematically and spatially with molecular emission intrinsically associated with W49A.  Based on comparisons with molecular data toward W49A, we suggest that the \\4\\/ \\HI\\/ component is directly associated with the northern part of the \\HII\\/ region ring, while the \\7\\/ \\HI\\/ component seems to originate in a lower density halo surrounding W49A.  We estimate that the W49A North core is significantly subvirial (2\\Te\\//\\We\\/$\\sim 0.2$), and that the total kinetic + magnetic energies amount to less than 1/3 of the total W49A North gravitational energy.  These magnetic field results suggest that W49A North is unstable to overall gravitational collapse in agreement with evidence that the halo is collapsing onto the W49A North ring of \\HII\\/ regions.  The majority of the \\HI\\/ column density toward W49B comes from Sagittarius Arm clouds along the line of sight at $\\sim 40$ \\kms\\/ and $\\sim 60$ \\kms\\/.  No significant magnetic fields were detected toward W49B.  Comparison of the spectral distribution of \\HI\\/ gas toward W49A and W49B suggests that evidence placing W49B 3 kpc closer to the sun (i.e.  at 8 kpc) than W49A is quite weak.  Although we cannot place W49B at the same distance as W49A, we find the morphology of a $\\sim 5$ \\kms\\/ \\HI\\/ component toward the southern edge of W49B suggestive of an interaction. ",
        "introduction": "The W49 complex is composed of one of the most luminous star forming regions in our galaxy (W49A) and a relatively young supernova remnant (W49B) $12\\arcmin$ to the east (see Fig.~\\ref{fig1}).  At high resolution, W49A is resolved into numerous compact \\HII\\/ regions (c.f.  Dreher et al.  1984; De Pree et al.  1997).  The \\HII\\/ regions to the north (W49A North) form a ``ring'' of small ($\\sim 5\\arcsec$) ultra-compact \\HII\\/ regions which appear to have formed during a period of triggered star formation (c.f.  Dickel \\& Goss 1990; also see Fig.~\\ref{fig2} for a magnified view of W49 North).  The \\HII\\/ region with the highest emission measure {\\bf G}\\footnote{In this paper we use the naming convention of De Pree et al.  (1997) for the W49A \\HII\\/ regions.}, is associated with a molecular outflow and some of the strongest \\hto\\/ masers in the galaxy (Scoville et al.  1986; Walker, Matsakis \\& Garcia-Barreto 1982).  To the SE of W49A North is W49A South which has a rim-brightened cometary shape and may have an edge-on interface with dense molecular gas to the west (Dickel \\& Goss 1990; also see Fig.~\\ref{fig2}).  Gwinn, Moran, \\& Reid (1992) recently estimated the distance to W49A to be 11.4 kpc from \\hto\\/ maser proper motions.  The giant molecular cloud associated with the W49A star forming region (as traced by \\tCO\\/) has a total mass of several $10^5$ M$_{\\odot}$ (Mufson \\& Liszt 1977).  \\placefigure{fig1}  Efforts to explain the complex morphology and kinematics of W49A have been abundant, and although progress has been made, thus far no single model can reproduce its observed properties.  Low excitation molecular emission lines (i.e.  \\tCO\\/($1-0$)) toward W49A are resolved into two components with velocities of \\4\\/ and 12 \\kms\\/ (Mufson \\& Liszt 1977).  Such observations led to the suggestion that there are two distinct clouds in the vicinity of W49A, and that their collision is responsible for triggering the observed star-formation (e.g. Mufson \\& Liszt 1977; Miyawaki, Hayashi, \\& Hasegawa 1986; Serabyn, G\\\"usten, \\& Schulz 1993).  However, observations of high excitation (i.e.  \\tCO\\/($7-6$)) and low abundance (i.e.  \\thCO\\/) lines have revealed the presence of colder, lower density gas at $\\sim 7$ \\kms\\/ which appears as self-absorption in more abundant species (Jaffe, Harris, \\& Genzel 1987; Phillips et al.  1981).  Moreover, inverse P Cygni profiles with absorption at $\\sim 12$ \\kms\\/ have been observed in high resolution ($\\sim 5\\arcsec$) observations of HCO$^+$, NH$_3$, and CS toward region {\\bf G} which indicate the presence of infall (Welch et al.  1987; Jackson \\& Kraemer 1994; Dickel et al.  1999).  These facts have been taken to imply that the infalling gas at $\\sim 12$ \\kms\\/ originates from a ``halo'' component at $\\sim 7$ \\kms\\/ and is falling onto the the ring of \\HII\\/ regions in W49A North.  Attempts to measure rotation in this gas or of the \\HII\\/ regions themselves have been largely unsuccessful or unreliable (Welch et al.  1987; Jackson \\& Kraemer 1994; De Pree, Mehringer, \\& Goss 1997; Dickel et al.  1999).  In addition to infall, there is also some evidence for the existence of more than one cloud or a velocity gradient in W49A (c.f. Dickel et al.  1999).  W49B is estimated to lie at a distance of $\\sim 8$ kpc (e.g.  Moffett \\& Reynolds 1994; also see \\S 4.2) and has one of the highest 1 GHz SNR surface brightnesses in the Galaxy.  It has a composite morphology with edge brightened radio synchrotron emission but centrally concentrated thermal X-ray emission (c.f Moffett \\& Reynolds 1994; Smith et al. 1985).  The spectral index of W49B is estimated to be $\\alpha\\sim -0.5$ ($S\\propto \\nu^{\\alpha}$) and the X-ray emission is thought to originate from a reverse shock in the remnant's interior (Green 1988; Kassim 1989; Moffett \\& Reynolds 1994).  The presence of such a reverse shock likely indicates that this remnant is young, with an age of $\\sim 6000$ years (Smith et al.  1985).  Moffett \\& Reynolds (1994; and references therein) find that the synchrotron emission of W49B has an unusually low fractional linear polarization, despite efforts to detect it at high frequencies and resolutions.  These authors note that the low degree of linear polarization observed toward W49B may be indicative of tangling in the magnetic field lines or Faraday depolarization within the SNR.  Diffuse atomic and molecular gas at LSR velocities of $\\sim 40$ and 60 \\kms\\/ which does not appear to be physically associated with W49 has also been observed toward W49A and W49B.  Some of the species observed in absorption include \\HI\\/, OH, H$_2$CO, HCO$^+$ and CS (Lockhart \\& Goss 1978, Bieging et al.  1982; Nyman 1984; Crutcher, Kaz\\`es, \\& Troland 1987; Miyawaki, Hasegawa, \\& Hayashi 1988; Greaves \\& Williams 1994).  The excitation requirements of these species suggest molecular hydrogen gas densities in the $10^2$ to $10^4$ cm$^{-3}$ range.  The $\\sim 40$ and 60 \\kms\\/ components are thought to arise from Sagittarius spiral arm clouds.  Since this line of sight passes through the Sagittarius arm twice, it is difficult to determine whether these clouds lie at their near or far distances (e.g.  Dame et al.  1986; Jacq, Despois, \\& Baudry 1988; also see \\S 4.2).  Given the wide range of physical conditions that can be sampled by a single observation toward W49 -- the complex kinematics of W49A, the SNR W49B, and diffuse gas at $\\sim 40$ and 60 \\kms\\/ coupled with the limited resolution of previous W49 \\HI\\/ observations ($2\\arcmin$; Lockhart \\& Goss 1978) we have performed VLA\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under a cooperative agreement by Associated Universities, Inc.}  \\HI\\/ Zeeman observations toward the W49 complex.  In particular, we hope to shed some light on the physical processes governing the star formation process(es) in W49A by imaging the magnetic field strengths and morphology in this intriguing source.  Previous VLA \\HI\\/ Zeeman magnetic field observations toward star forming regions like W3 (Roberts et al.  1993) and M17 (Brogan et al.  1999) have revealed line of sight field strengths of several hundred \\muG\\/ in gas intrinsic to these sources, as well as, complex field morphologies.  The results of our VLA \\HI\\/ Zeeman absorption observations toward the W49 complex are presented below and are organized in the following order:  details of the data reduction and observing parameters are given in \\S 2; discussion of the 21 cm W49A and W49B continuum morphology can be found in \\S 3.1, while the properties of the \\HI\\/ absorption toward W49A and W49B are described in \\S 3.2; our \\HI\\/ Zeeman magnetic field results are presented in \\S 3.3; and finally, the implications of these data are discussed in \\S 4. ",
        "conclusions": "The \\HI\\/ gas in the velocity range ($-5$ to 25 \\kms\\/) shows good agreement both kinematically and spatially with molecular emission intrinsically associated with W49A (see Figs.~\\ref{fig4} and ~\\ref{fig5}).  Therefore, the \\HI\\/ gas in this velocity range is likely to originate near the W49A star-forming complex.  Significant line of sight magnetic fields toward W49A were detected in two different \\HI\\/ velocity components.  Some of the properties of these detections are:  (1) \\Blos\\/ of 60 to 300 \\muG\\/ were detected toward W49A at velocities of \\4\\/ and \\7\\/; (2) {\\em The \\4\\/ \\Blos\\/ is negative, while the \\7\\/ is positive} (negative \\Blos\\/ indicates that the line of sight field points toward the observer); and (3) The \\Blos\\/ values measured toward W49A show a significant {\\em increase in field strength with higher resolution} especially for the \\4\\/ \\HI\\/ component (see Figs.~\\ref{fig10}, ~\\ref{fig11}, ~\\ref{fig12}, ~\\ref{fig13}, and ~\\ref{fig14}).  Based on comparisons of the \\4\\/ and \\7\\/ \\Blos\\/ detection regions and velocities with molecular data toward W49A, it seems likely that the \\4\\/ \\HI\\/ component is directly associated with the northern part of the \\HII\\/ region ring, while the \\7\\/ \\HI\\/ component seems to originate in a lower density halo surrounding W49A.  From previous molecular line and dust observations of W49A, we estimate that the W49A North core ($\\sim 1\\arcmin$ in extent) is significantly subvirial, with 2\\Te\\//\\We\\/$\\sim 0.2$.  In addition, it appears that the total magnetic energy (assuming ${\\mid\\bf\\vec{B}}\\mid\\sim 600$ \\muG\\/ and \\Mw\\/$\\sim$ \\Te\\/) only contributes 0.1\\We\\/.  Adding the different energy contributions, we estimate that the total kinetic + magnetic energies only amount to less than about 1/3 of the total gravitational energy.  Indeed, we find that ${\\mid\\bf\\vec{B}}\\mid\\sim 5$ mG is needed for there to be overall equilibrium in W49A North (ignoring external pressure).  Therefore, while it must be emphasized that the parameters used in these estimates are fairly uncertain, these new magnetic field results suggest that W49A North is unstable to overall gravitational collapse as implied by evidence that the halo is collapsing onto the W49A North ring of \\HII\\/ regions (e.g.  Dickel et al.  1999).  There are four main \\HI\\/ velocity groups detected toward W49B with velocities of $\\sim 60$, $\\sim 40$, $\\sim 20$, and $\\sim 5$ \\kms\\/.  The majority of the \\HI\\/ column density toward W49B comes from the $\\sim 60$ and $\\sim 40$ \\kms\\/ \\HI\\/ components which arise from Sagittarius Arm clouds along the line of sight.  No significant \\Blos\\/ fields were detected toward W49B.  Comparison of the spectral distribution of \\HI\\/ gas toward W49A and W49B (Fig.~\\ref{fig16}) suggests that there is no clear evidence that W49B is 3 kpc closer to the sun than W49A (known to be at 11.4 kpc from water maser proper motions; Gwinn et al.  1992).  Although we cannot place W49B at the same distance as W49A, we find the morphology of the $\\sim 5$ \\kms\\/ \\HI\\/ gas toward the southern edge of W49B suggestive of an interaction. Indeed, gas at similar velocities is known to be intrinsically associated with W49A."
    },
    "0011/gr-qc0011096_arXiv.txt": {
        "abstract": "s{We present tests and results of a new axisymmetric, fully general relativistic code capable of solving the coupled Einstein-matter system for a perfect fluid matter field. Our implementation is based on the Bondi metric, by which the spacetime is foliated with a family of outgoing light cones. We use high-resolution shock-capturing schemes to solve the fluid equations. The code can accurately maintain long-term stability of a spherically symmetric, relativistic, polytropic equilibrium model of a neutron star. In axisymmetry, we demonstrate global energy conservation of a perturbed neutron star in a compactified spacetime, for which the total energy radiated away by gravitational waves corresponds to a significant fraction of the Bondi mass.} ",
        "introduction": " ",
        "conclusions": "We have developed an axisymmetric, fully general relativistic characteristic code with a perfect fluid matter field. We have checked convergence in spacetimes containing a neutron star.  For a spherical, relativistic equilibrium model of a polytropic neutron star, we have shown long-term stability for hundreds of light-crossing times. In axisymmetry, we analyzed convergence in strongly perturbed neutron stars. Compactifying the spacetime, we could show global energy conservation, in the sense that the amount of energy lost by the system is radiated away by gravitational waves."
    },
    "0011/astro-ph0011018_arXiv.txt": {
        "abstract": "{ The RS\\,CVn binary system HR\\,1099 was extensively observed by the {\\it XMM-Newton} observatory in February 2000 as its first-light target. A total of 570\\,ks of exposure time was accumulated with the Reflection Grating Spectrometers (RGS). The integrated X-ray spectrum between 5--35\\AA\\ is of unprecedented quality and shows numerous features attributed to transitions of the elements C, N, O, Ne, Mg, Si, S, Fe, and Ni. We perform an in-depth study of the elemental composition of the average corona of this system, and find that the elemental abundances  strongly depend on the first ionisation potential (FIP) of the elements. But different from the solar coronal case, we find an {\\it inverse FIP effect}, i.e., the abundances (relative to oxygen) {\\it increase with increasing FIP}. Possible scenarios, e.g., selective enrichment due to Ne-rich flare-like events, are discussed. ",
        "introduction": "The outer atmospheres of stars are of fundamental importance to the chemical enrichment of interstellar space through stellar winds and the ejection of energetic particles (Meyer 1985). Interestingly, the elemental composition of the hot magnetic solar corona and the solar wind are markedly different from the Sun's surface composition, which suggests that some fractionation process in the denser layers of the solar atmosphere must selectively enrich the corona with particular elements. Various observations of the solar corona, the solar wind, solar energetic particles, and cosmic rays have established that elements with a first ionisation potential (FIP) below 10\\,eV are overabundant by factors of up to 4 with respect to the solar photosphere, unlike elements with a high FIP that maintain their photospheric abundance (Feldman 1992, see also Meyer 1985). This is usually referred to as the ``FIP effect\". However, the physics of this fractionation mechanism is poorly understood. Comparative abundance studies of stars other than the Sun may shed light on the underlying physics.  \\begin{figure*}[ht] \\resizebox{12cm}{!}{\\includegraphics[angle=0, clip, bb=43 41 716 525]{XMM40_f1.eps}} \\hfill \\parbox[b]{55mm}{ \\caption{First order spectrum of HR\\,1099 as measured by the RGS instruments on board XMM-Newton. The most prominent spectral features are indicated just above the spectrum by the corresponding emitting ion.} \\vspace{1cm} } \\end{figure*}  The X-ray regime contains a rich forest of atomic spectral lines emitted by many chemical species at ionisation stages corresponding to plasma temperatures between $\\sim$1 and 30 Million K.  So far, however, attempts to determine the composition of stellar coronae other than the Sun's from X-ray observations (Antunes et al.\\ 1994, White 1996, Drake 1996) have produced results that are considered to be ambiguous, because the low spectral resolution available from even the best nondispersive detectors has not permitted the separation of atomic emission lines arising from different elements. Often, subphotospheric abundances were claimed. Studies of stellar coronal abundances with the EUVE satellite have revealed solar-like FIP effects in some, but not all of the stars studied (Drake et al.\\ 1995, 1997). The high-resolution grating spectrometers on board the recently launched American Chandra and European XMM-Newton satellites now provide a new opportunity to unambiguously determine both the elemental abundances and the thermal structure of stellar coronae through explicit measurement of emission line and continuum intensities in the X-ray band. We are using grating spectroscopy on XMM-Newton to perform such analyses for the first time. The target chosen for the present observation is HR\\,1099, one of the X-ray brightest and most magnetically active stellar systems, which was observed early in the programme. HR\\,1099 has been the target of previous EUVE and low-resolution X-ray spectroscopic investigations (e.g.\\ Pasquini et al.\\ 1989, Griffith \\&\\ Jordan 1998). Ayres et al.\\ (2001) qualitatively compare the Chandra HETG spectra of HR\\,1099 and Capella, the latter being a somewhat cooler and less active coronal source. A companion paper (Audard et al.\\ 2001) to the present work discusses spectral variability in HR\\,1099, based on selected XMM-Newton RGS data sets. ",
        "conclusions": "The Ne/O ratio found here for the HR\\,1099 coronae is unusually high compared with the solar photospheric value, while the N/O ratio is also enhanced and the Mg/O ratio is close to solar photospheric. The abundance ratios for the remaining elements are well below their solar photospheric values. Compared with the respective ratios of the solar corona, the contrast is even sharper. In fact, HR\\,1099 reveals systematically a reversed FIP effect, which is illustrated in Fig.~4; the coronal enrichment increases monotonically with the FIP. The only exception may be Mg (low FIP). Ne is strongly enhanced, by a factor of 3.8 compared to the solar photosphere. This enhancement is reminiscent of a similar anomaly observed in a subset of solar flares (Schmelz 1993, Murphy et al.\\ 1991). Interestingly, flares are thought to be of great importance to the heating of apparently quiet coronae of active stars (Doyle \\&\\ Butler 1985, G\\\"udel 1994, Audard et al.\\ 2000) and have been suggested to be pivotal for the elemental fractionation process as well (Wang 1996, G\\\"udel et al.\\ 1999, Mewe et al.\\ 1997). Continuous mixing of photospheric and coronal material in low-lying loops through frequent flares could suppress the solar-like FIP effect (Feldman \\&\\ Widing 1990) in these stars altogether.  It is generally believed that the fractionation is the result of diffusive processes across magnetic field lines somewhere in the chromosphere, where the ionisation fraction differs substantially among elements of high or low FIP (von Steiger \\&\\ Geiss 1989). Low FIP elements can then be accelerated into the quiescent corona, leaving an excess of high FIP elements in the chromosphere (Schmelz 1993). If this chromospheric layer can be accessed by flares (Schmelz 1993), then a high-FIP-rich plasma may be heated and lifted into dense coronal X-ray loops giving rise to the reversed coronal FIP effect observed in HR\\,1099. In any case, it is clear that the simple mechanism involving steady-state diffusion, which has been proposed for the quiet solar corona, does not operate in more active stars such as HR\\,1099. Since cosmic rays do show a low FIP enhancement, it is unlikely that such active stars can be major contributors to the cosmic ray flux. The wide survey of stellar coronae planned with the RGS on XMM-Newton will provide more general answers to the questions raised by these observations."
    },
    "0011/astro-ph0011368_arXiv.txt": {
        "abstract": "We report on long-term monitoring of the anomalous X-ray pulsar (AXP) \\psr\\ using the {\\it Rossi X-ray Timing Explorer}.  The timing behavior of this pulsar is different from that of other AXPs being monitored with \\xte. In particular, we show that the pulsar shows significant deviations from simple spin-down such that phase-coherent timing has not been possible over time spans longer than a few months.  We find that the deviations from simple spin down are not consistent with single ``glitch'' type events, nor are they consistent with radiative precession.  We show that in spite of the rotational irregularities, the pulsar exhibits neither pulse profile changes nor large pulsed flux variations.  We discuss the implications of our results for AXP models. We suggest that \\psr\\ may be a transition object between the soft gamma-ray repeater and AXP populations, and the AXP most likely to one day undergo an outburst. ",
        "introduction": "The nature of anomalous X-ray pulsars (AXPs) has been a mystery since the discovery of the first example some 20 years ago. Although it is clear that they are young neutron stars, it is not clear why they are observable: their X-ray emission defies conventional explanations that rely on either accretion or rotation power as the radiation energy source.  Although only five AXPs are known, their origin is likely to be of great importance to our understanding of the fate of massive stars and the basic properties of the young neutron star population.  AXP observational properties can be summarized as follows (e.g. \\cite{ms95}): (i) they have spin periods in the narrow range $\\sim$6--12~s; (ii) they exhibit regular spin-down on time scales of $10^3$--$10^5$~yr; (iii) they show no evidence for Doppler shifts due to binary motion; (iv) they lie in the Galactic plane, implying youth; (v) two appear to be located at the geometric centers of supernova remnants (SNRs), also indicating youth; (vi) all have much softer X-ray spectra than accretion-powered pulsars; (vii) they have no known counterparts at any other energy range, including radio or optical (but see below); and (viii) their X-ray luminosities greatly exceed their inferred spin-down powers, hence rotation power is implausible. For an excellent recent review of these objects, see Mereghetti (1999).  Suggested models for AXPs fall into two broad categories. One model proposes that AXPs are young, isolated, highly magnetized neutron stars or ``magnetars'' (\\cite{dt92a}).  In this case, AXP spin-down is due either to magnetic dipole radiation as in radio pulsars, which implies large surface magnetic fields, $10^{14}-10^{15}$~G, or it could be due to angular momentum loss from Alfv\\'en wave emission following crust cracking due to stresses caused by an even larger magnetic field (\\cite{tb98,tdw+00,dun00}).  Similar fields are inferred independently in the soft gamma repeaters (SGRs) which, aside from their soft gamma-ray outbursts, exhibit AXP-like X-ray pulsations in quiescence (\\cite{kf93,td95,kds+98}).  In this model, the AXP X-ray luminosities could come from the decay of the large magnetic field (\\cite{td96a}) or from enhanced thermal emission from the young cooling neutron star (\\cite{hh97}).  The magnetar model does not presently explain all AXP properties, however; for example it is unclear why all the observed rotation rates should fall in so narrow a range.  The second model of AXP emission is that they are powered by accretion.  Conventionally, accreting material would come from a companion (\\cite{ms95}), however the absence of Doppler shifts allows only very low mass and compact companions unless all known systems are seen fortuitously face-on (\\cite{mis98,wdf+99}).  The companion may have been disrupted leaving behind a fossil accretion disk (\\cite{vtv95,cso+95}), although neither scenario is consistent with the apparent youth of these objects.  An alternative is that AXPs are neutron stars with accretion disks of material leftover from the supernova explosion (\\cite{chn00,alp99,mlrh01,ch00}). However, very constraining optical observations of one AXP, 1E~2259+586, detect no evidence for any companion, nor evidence for an accretion disk (\\cite{hkvk00}). A recent detection of a possible very faint optical counterpart to the AXP 4U~0142+61 also renders the presence of an accretion disk problematic (\\cite{hvk00}).  One way to distinguish between these classes of models may be through timing observations.  In the magnetar model for AXPs, relatively smooth spin-down should be expected, punctuated by occasional abrupt spin-up or spin-down events or ``glitches,'' as well as low-level, long-time-scale deviations from simple spin-down, or ``timing noise.'' Both phenomena are well known among young radio pulsars (e.g. \\cite{lyn96}), although their physical origins in magnetars may be considerably different given the much larger inferred magnetic field. Indeed large amounts of timing noise have been seen in SGRs (e.g. \\cite{wkf+00}).  However no extended spin-up should be seen in the magnetar model.  Further, Melatos (1999) \\nocite{mel99} suggested that deformation of the neutron star due to a large magnetic field would result in radiative precession, which would manifest itself observationally as a long-term periodicity in timing data.  On the other hand, accretion power is usually associated with much noisier timing behavior, which can be correlated with spectral, luminosity, and pulse morphology changes.  In addition, some accreting binary systems undergo extended ($\\sim$years) episodes of spin-up, although these generally seem to alternate with long intervals of spin-down as well (\\cite{bcc+97}).  Until recently, monitoring of AXPs produced only measured frequencies at individual epochs separated by many months or years.  This method of timing is not sensitive to spin irregularities on a wide range of interesting time scales.  Such ``incoherent'' timing can offer no insight into the spin behavior of the source between observations, nor to the abruptness of apparent changes in spin-down rate.  For example, although 1E 2259$+$586 and \\psr\\ have the longest recorded timing histories of all AXPs (e.g. \\cite{cm97,bsss98,pkdn00}), the histories are poorly sampled and observed apparent deviations from simple spin-down have been given very different interpretations (e.g. \\cite{mel99,hh99,bss+00}).  Kaspi, Chakrabarty \\& Steinberger (1999) [hereafter KCS99], \\nocite{kcs99} making use of the scheduling flexibility of the {\\it Rossi X-ray Timing Explorer}, showed that ``phase coherent timing,'' that is, long-term monitoring that permits absolute pulse numbering between observations (a technique that has been used for radio pulsars for many years), can successfully be applied to at least two AXPs. This implies that AXPs can be extremely stable rotators: KCS99 showed that for AXPs 1E~2259+586 and RXS J170849.0$-$400910, deviations from a simple linear spin-down model described the rotational behavior to within a few percent of the pulse period over years.  Such stability, together with phase-coherent timing, allowed Kaspi, Lackey, and Chakrabarty (2000) \\nocite{klc00} to detect a sudden spin-up event having all the hallmarks of a classical radio pulsar glitch in RXS~J170849.0$-$400910.  Whether such glitches are ubiquitous in AXPs, and how their amplitude distributions and frequencies compare to those well studied in radio pulsars, are unknown.  Similarly, whether great rotational stability during glitch-free intervals is common to all AXPs is also unknown, as is the origin of deviations from simple spin-down that have been seen in the long-term timing histories of 1E~2259+586 and \\psr.  \\psr\\ is a 6.4~s X-ray pulsar discovered in {\\it Einstein} observations of the Carina nebula (\\cite{scs86}).  It has been shown to have a soft X-ray spectrum (e.g. \\cite{opmi98} and references therein), like other AXPs.  It exhibits no evidence for any binary companion, as no Doppler shifts of the pulse period are seen (\\cite{mis98}), and no optical counterpart to a limiting magnitude of $m_V \\sim 20$ has been detected (\\cite{mcb92}).  Occasional monitoring observations over more than 20 years show that the pulsar is spinning down, though significant deviations from a simple spin-down model have been noted (\\cite{mer95,cm97,opmi98,pkdn00,bss+00}). The paucity of data thus far has made it impossible to unambiguously identify the origin of the deviations.   Here we report on our monthly {\\it Rossi X-Ray Timing Explorer} (\\xte) monitoring of \\psr\\ in which we have attempted phase-coherent timing like that achieved for 1E~2259+586 and RXS J170849.0$-$400910.  We show that the spin behavior of \\psr\\ is qualitatively and quantitatively different from those of the other two AXPs.  In particular, in our observations, \\psr\\ has exhibited rotation that is much less stable than those of the other AXPs timed as part of our \\xte\\ program.  We also show that \\psr\\ occasionally suffers rotational anomalies on short time scales that cannot be described by simple spin-down laws, classical glitches, or radiative precession.  Further, we use the same \\xte\\ observations to search for pulse profile changes and/or pulsed flux variations, and to check for correlations with timing behavior. We find no such correlations.  In \\S\\ref{sec:obs} we describe our observations, in \\S\\ref{sec:results} we describe our analysis procedures and results, and we discuss the implications of our results in \\S\\ref{sec:disc}. ",
        "conclusions": "Long-term \\xte\\ monitoring of the AXP \\psr\\ has shown it to be a noisier rotator than the other AXPs monitored by \\xte\\ (KCS99). Indeed, phase-coherent timing over spans longer than a few months has not been possible for \\psr.  Furthermore, we have observed deviations from simple spin-down that are inconsistent with a single glitch event, and inconsistent with being quasi-periodic as predicted for radiative precession.  Only denser sampling will help determine the nature of the spin-down irregularities in this source.  Our \\xte\\ observations also show that the pulsar has exhibited no pulse profile variations nor large pulsed flux variations.  The latter is inconsistent with the simplest predictions of accretion theory, although deep optical/IR upper limits are the best hope for conclusively ruling out an accretion origin of the X-ray emission. We note that the spin and spectral properties of \\psr\\ are, among AXPs, the most reminiscent of the SGRs. We therefore suggest that this pulsar may be a transition object between the two populations, and the AXP most likely to one day undergo an outburst."
    },
    "0011/astro-ph0011474_arXiv.txt": {
        "abstract": "s{ We perform a detailed comparison of the \\iras\\ \\pscz\\ and \\jy\\  spherical harmonic coefficients of the predicted velocity fields in redshift space. The monopole terms predicted from the two surveys show some differences. Faint galaxies are responsible for this mismatch that disappears when extracting a \\pscz\\ subsample of galaxies with fluxes larger than \\jy. The analysis of \\pscz\\ dipole components confirms the same inconsistencies found by Davis, Nusser and Willick \\cite{Davis:1996} between the \\iras\\ \\jy\\ gravity field and \\markiii\\ peculiar velocities. Shot-noise, which is greatly reduced in our \\pscz\\ gravity field, cannot be responsible for the observed mismatch. } ",
        "introduction": "Nusser \\& Davis (NS)\\cite{Nusser:1994} show that in linear gravitational instability (GI) theory the peculiar velocity field in redshift space is irrotational and thus can be expressed in terms of a potential: ${\\vec v}=-\\nabla \\Phi(\\vec s)$. The angular dependencies of the potential velocity field and the galaxy overdensity field [both measured in redshift space and expanded in spherical harmonics, $\\Phi_{lm}(s)$ and $\\hat{\\delta}_{lm}(s)$, respectively] are related by a modified Poisson equation: \\begin{equation} {{1}\\over{s^2}}{{d}\\over{ds}} \\left( s^2 {{d\\Phi_{lm}}\\over{ds}} \\right) -{{1}\\over{1+\\beta}} {{l(l+1)\\Phi_{lm}}\\over{s^2}}= {{\\beta}\\over{1+\\beta}} \\left( \\hat{\\delta}_{lm}- {{1}\\over{s}}{{d\\ln{\\phi}}\\over{d \\ln{s}}} {{d \\Phi_{lm}}\\over{ds}} \\right), \\label{eq:ndchapter5} \\end{equation} where $\\phi(s)$ is the selection function and $\\beta \\equiv \\Omega ^{0.6}/b$, where $\\Omega$ is the matter density and $b$ the bias of the galaxy distribution. To solve this differential equation we first compute the density field on an angular grid using cells of equal solid angle and 52 bins in redshift out to $s=18\\,000$~kms$^{-1}$. \\begin{equation} 1+{\\hat \\delta}_j({\\vec s}_n)=\\frac{1}{(2\\pi)^{3/2}\\sigma_{1.2 n}^3}\\sum_i^{N_j} \\frac{1}{\\phi(s_i)}\\exp\\left [ -\\frac{({\\vec s}_n-{\\vec s}_i)^2}{2\\sigma_{1.2 n}^2} \\right ] \\label{eq:smoothing5} \\end{equation} where the sum is over all the galaxies within the catalogue $j$, $N_j$.  The Gaussian smoothing width for the cell $n$ at redshift $s_n$, $\\sigma_{1.2 n}$, is given by $\\sigma_{1.2 n} = \\mbox{max}\\{100, [\\bar{n}_{1.2}\\phi_{1.2}(s_n)]^{-1/3}\\}$\\ km s$^{-1}~$, where $\\bar{n}_{1.2}$~and $\\phi_{1.2}$~are the 1.2 Jy mean number density and selection function, respectively. ",
        "conclusions": "\\label{section:conclusions5} \\begin{figure} \\begin{center} \\scalebox{0.70}{\\includegraphics[37,472][538,755]{lteodoro1.eps}} \\caption{Monopole coefficient, $u_{00}(s)$, estimated from \\pscz~and 1.2-Jy catalogues for $\\beta =0.6$. On the two top panels dashed (thick) lines represent 1.2-Jy (\\pscz) monopole mode of the velocity field.  a) The hatched and shaded regions indicate the 1.2 Jy and \\pscz\\ 1--$\\sigma$\\ uncertainties, respectively. b) Dotted, dot-dashed and dashed lines represent \\pscz$_{0.7}$, \\jy\\  and \\pscz$_{\\mbox{1.2}}$, respectively. } \\label{fig:irasmonopole} \\end{center} \\end{figure}  We compare the line of sight peculiar velocities relative to the Local Group, $u ({\\vec s})\\equiv [ {\\vec v}({\\vec v}) -{\\vec v}_{LG} ] \\cdot {\\hat s}$, obtained from the \\iras\\ \\pscz\\ and \\jy\\ redshift surveys gravity fields. The comparison is performed in redshift-space.  To perform the decomposition in spherical harmonics coefficients we apply $\\beta=0.6$.  In Fig.~\\ref{fig:irasmonopole} we display the monopole of the velocity field, $u_{00}$. In the left panel, the estimate of the velocity monopole in the \\jy\\ (dashed line) is systematically larger than the \\pscz\\ one (continuous line) in the range $1\\,000 < s < 8\\,000 $~km s$^{-1}$.  In Fig.~\\ref{fig:irasdipoles} we show the three velocity dipole components of the two surveys along with the total amplitude (bottom left panel). The various dipole components exhibit good agreement, specially $u_{1,-1}$ (top right panel). Davis, Nusser and Willick (DNW)\\cite{Davis:1996} find systematic discrepancies between the \\markiii\\  and \\iras\\ \\jy--predicted flow fields, particularly a difference in the  $m=-1$ component of the dipole. Estimating the velocity field from the \\iras\\ \\pscz\\ catalogue allows us to understand the influence of sparse-sampling on the \\jy\\ inferred fields. Since both dipoles are so similar, we might  conclude that the discrepancies with the \\markiii\\ velocities do not disappear when the shot noise is reduced by a factor $\\approx  \\sqrt 3$, like in the \\pscz\\  velocity field.   Willick \\etal\\ applied their VELMOD machinery to compare the observed velocities of the \\markiii\\ galaxies with those predicted by the \\iras\\ \\jy\\ model within $ s = 3\\,000$\\  $\\kms$. They found that the residuals are well modeled by a quadrupole with an amplitude of $\\sim 3 \\%$\\ that of the Hubble flow.  In Fig. 3, the quadrupole components inferred from both \\iras\\ catalogues agree quite well, except that the $m=0,2$ magnitudes are somewhat larger in \\jy, therefore a quadrupole mismatch between \\markiii\\ and the \\pscz\\ velocities is also to be expected. Note that when all the multipoles are considered ({\\it{i.e.}}\\ when performing a full v-v comparison) the \\pscz\\  and \\jy\\  gravity fields look fully consistent (B99).\\cite{Branchini:1999}  In all plots, shaded and hatched regions represent 1-$\\sigma$ uncertainties. These incorporate sources of errors in the velocity field: {\\it i)} the shot-noise error due to sparse-sampling of the underlying density and velocity fields; {\\it ii)} systematic and random errors caused by a less than perfect method of reconstruction. The uncertainties are evaluated summing in quadrature the errors due to shot-noise and filling-in procedure.  We compute the shot-noise error by generating 100 bootstrap realizations of the observed distributions of \\iras\\ galaxies and computing the velocity fields from these realizations. To quantify random and systematic errors caused by the filling-in procedure we use a suite of 20 mock-catalogues that mimic the main properties of the \\pscz\\ and \\jy\\ redshift surveys. A complete description of the  prescription followed to quantify the velocity uncertainties can be found in Teodoro, Branchini and Frenk \\cite{Teodoro:2000}.   Where does the discrepancy between the monopoles of the \\pscz\\ and \\jy\\ surveys come from? As shown in the plots, the difference is larger than that expected from the shot--noise (included in the error budget). Tadros {\\it et al.}\\cite{Tadros:1999} have suggested that the \\pscz\\ catalogue may be incomplete for fluxes $\\le 0.7$~Jy. If true, then we would expect that the velocity monopole for the \\pscz\\ with a flux cut at 0.7 Jy (\\pscz$_{0.7}$) would be in good agreement with the 1.2-Jy survey. The dotted line ($u_{00,\\mbox{\\pscz}_{0.7}}$) in the top panel of Fig. \\ref{fig:irasmonopole} shows shows that objects fainter than 0.7 Jy are only partially responsible for the monopole  mismatch that therefore cannot be ascribed to incompleteness at low fluxes however that this is not case. It is only in cutting the \\pscz\\ catalogue at a flux level of 1.2 Jy that, as expected, the discrepancy disappears. This is clearly seen in the right panel of Fig. \\ref{fig:irasmonopole} in which the dashed line indicates  \\pscz$_{1.2}$ velocity monopole.  We conclude that a better sampled \\iras\\ catalogue cannot resolve alone the mismatches between the \\markiii\\ peculiar velocity field and the \\iras\\ predicted gravity field.    \\begin{figure} \\begin{center} \\scalebox{0.55}{\\includegraphics[24,268][526,690]{lteodoro2.eps}} \\caption{ The dipole coefficients, $u_{1m}(s)$\\{$m=-1,0,1$\\}, inferred from \\pscz\\ and \\jy\\ for $\\beta =0.6$. The solid and dashed lines in the various panels represent \\pscz\\ and \\jy\\ dipoles, respectively. The $GX$, $GY$, $GZ$~panels show the three Galactic components of the dipole and the bottom left is their quadrature sum. Hatched and shaded regions indicate the \\jy\\ and \\pscz\\ 1--$\\sigma$\\ uncertainty, respectively.} \\label{fig:irasdipoles} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\scalebox{0.55}{\\includegraphics[24,27][526,766]{lteodoro3.eps}} \\caption{ The quadrupole coefficients, $u_{1m}(s)$\\{$m=-2,-1,0,1,2$\\}, inferred from \\pscz\\ and 1.2-Jy for $\\beta =0.6$. The solid and dashed lines in the various panels represent \\pscz\\ and 1.2-Jy quadrupoles, respectively. The $m = 0,\\pm 1, \\pm 2 $\\ panels show the five Galactic components of the quadrupole and the bottom left is their quadrature sum. Hatched and shaded regions indicate the 1.2-Jy and \\pscz~1--$\\sigma$\\ uncertainty, respectively.} \\label{fig:irasquadrupole} \\end{center} \\end{figure}"
    },
    "0011/astro-ph0011197_arXiv.txt": {
        "abstract": "We present an alternative scenario for cosmic structure formation where initial fluctuations are due to Kalb-Ramond axions produced during a pre-big bang phase of inflation. We investigate whether this scenario, where the fluctuations are induced by seeds and therefore are of isocurvature nature, can be brought in agreement with present observations by a suitable choice of cosmological parameters. We also discuss several observational signatures which can distinguish axion seeds from standard inflationary models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011105_arXiv.txt": {
        "abstract": "The emission-line flux ratio of [O {\\sc iii}]$\\lambda$4363/[O {\\sc iii}]$\\lambda$5007 ($R_{\\rm O III}$) is a useful diagnostic for the ionization mechanism and physical properties of emission-line regions in active galactic nuclei (AGNs). However, it is known that simple photoionization models underpredict the [O {\\sc iii}]$\\lambda$4363 intensity, being inconsistent with observations. In this paper, we report on several pieces of evidence that a large fraction of the [O {\\sc iii}]$\\lambda$4363 emission arises from the dense gas obscured by putative tori: (1) the visibility of high-$R_{\\rm O III}$ regions is correlated to that of broad-line regions, (2) higher-$R_{\\rm O III}$ objects show hotter mid-infrared colors, (3) higher-$R_{\\rm O III}$ objects show stronger highly-ionized emission lines such as [Fe {\\sc vii}]$\\lambda$6087 and [Fe {\\sc x}]$\\lambda$6374, and (4) higher-$R_{\\rm O III}$ objects have broader line width of [O {\\sc iii}]$\\lambda$4363 normalized by that of [O {\\sc iii}]$\\lambda$5007. To estimate how such a dense component contributes to the total emission-line flux, dual-component photoionization model calculations are performed. It is shown that the observed values of $R_{\\rm O III}$ of type 1 AGNs may be explained by introducing a 5\\% -- 20\\% contribution from the dense component while those of type 2 AGNs may be explained by introducing a 0\\% -- 2\\% contribution. We also discuss the [O {\\sc iii}]$\\lambda$4363 emitting regions in LINERs in the framework of our dual-component model. ",
        "introduction": "It has been often considered that emission-line regions around active galactic nuclei (AGNs) are photoionized by the nonthermal continuum radiation from central engines (e.g., Davidson 1977; Yee 1980; Kwan \\& Krolik 1981; Shuder 1981; Cohen 1983; Cruz-Gonz\\'alez et al. 1991; Osterbrock 1993; Evans et al. 1999). However, this photoionization scenario has sometimes been confronted with several serious problems.  One of such problems is that any single-zone photoionization models underpredict the [O {\\sc iii}]$\\lambda$4363/[O {\\sc iii}]$\\lambda$5007 intensity ratio, $R_{\\rm O III}$ (e.g., Koski \\& Osterbrock 1976; Ferland \\& Netzer 1983; Filippenko \\& Halpern 1984; Viegas-Aldrovandi \\& Gruenwald 1988). The reason for the underprediction of $R_{\\rm O III}$ is thought to be that photoionization of gas in optically thick condition is hard to accomplish electron temperatures above a few $\\times 10^4$ K if density of the gas is typical in narrow-line regions (NLRs). In order to solve this problem, many studies have been carried out. Such attempts can be roughly divided into the following two categories. One is based on the idea that a high density component may contribute to achieve the observed high $R_{\\rm O III}$ (e.g., Baldwin 1975; Osterbrock, Koski, \\& Phillips 1976; Filippenko \\& Halpern 1984; Filippenko 1985). This idea is attributed to the fact that the critical density ($n_{\\rm cr}$) of the [O {\\sc iii}]$\\lambda$4363 emission (3.3 $\\times 10^7$ cm$^{-3}$) is higher than that of the [O {\\sc iii}]$\\lambda$5007 emission (7.0 $\\times 10^5$ cm$^{-3}$); which leads to high $R_{\\rm O III}$ when the gas density is higher than $\\sim 10^6$ cm$^{-3}$. The other idea is to introduce high temperature regions whose temperature is more than a few times 10$^4$ K. To achieve such high temperatures, either shock-heated regions (e.g., Koski \\& Osterbrock 1976; Heckman 1980; Dopita \\& Sutherland 1995) or optically-thin components (e.g., Wilson, Binette, \\& Storchi-Bergmann 1997) is required.  In addition, some previous studies reported that the values of $R_{\\rm O III}$ depend on the AGN type; i.e., type 1 Seyfert nuclei (S1s) exhibit higher $R_{\\rm O III}$ than type 2 Seyfert nuclei (S2s) (e.g., Osterbrock et al. 1976; Heckman \\& Balick 1979; Shuder \\& Osterbrock 1981; Cohen 1983). This tendency seems to suggest that a part of the [O {\\sc iii}]$\\lambda$4363 emission arises from the regions which are obscured only in S2s by any materials, such as dusty tori. However, this tendency may be interpreted by the intrinsic (that is, not due to obscuration effects) difference of NLR properties, such as the size (Schmitt \\& Kinney 1996; Kraemer et al. 1998) or the ionization states (Schmitt 1998).  To investigate the reason why S1s show high $R_{\\rm O III}$ than S2s seems useful to explore where and how the [O {\\sc iii}]$\\lambda$4363 emission is radiated. Moreover, these may lead to the new solution to the underprediction problem of $R_{\\rm O III}$. Therefore, in this paper, we present how the observed values of $R_{\\rm O III}$ are different among various types of Seyferts based on a large sample of Seyferts compiled from the literature. Then, we compare $R_{\\rm O III}$ with various parameters and discuss the nature of the [O {\\sc iii}]$\\lambda$4363 emitting regions in AGNs. ",
        "conclusions": "\\subsection{Where is the [O {\\sc iii}]$\\lambda$4363 Emitting Region in AGNs?}  \\begin{figure*} \\epsscale{0.9} \\plotone{fig10.ps} \\caption{ The frequency distributions of the [O {\\sc iii}]$\\lambda$5007 luminosity (left) and the [O {\\sc iii}]$\\lambda$4363 luminosity (right) for the NLS1s, the BLS1s, the S1.5s, and the S2$_{\\rm total}$s. \\label{fig10}} \\end{figure*}  \\begin{deluxetable}{lcccc} \\tablenum{6} \\tablecaption{The Results of the KS Test Concerning the $L_{\\rm [O III]\\lambda 5007}$ and $L_{\\rm [O III]\\lambda 4363}$ \\label{tbl-6}} \\tablewidth{0pt} \\tablehead{ \\colhead{Type} & \\colhead{NLS1} & \\colhead{BLS1} & \\colhead{S1.5} & \\colhead{S2$_{\\rm total}$} } \\startdata \\cutinhead{$L_{\\rm [O III]\\lambda 5007}$} NLS1             & \\nodata & 4.291 $\\times 10^{-3}$ & 4.394 $\\times 10^{-4}$ & 4.445 $\\times 10^{-2}$ \\nl BLS1             & \\nodata & \\nodata                & 9.312 $\\times 10^{-1}$ & 8.359 $\\times 10^{-2}$ \\nl S1.5             & \\nodata & \\nodata                & \\nodata                & 1.225 $\\times 10^{-2}$ \\nl S2$_{\\rm total}$ & \\nodata & \\nodata                & \\nodata                & \\nodata                \\nl \\cutinhead{$L_{\\rm [O III]\\lambda 4363}$} NLS1             & \\nodata & 2.519 $\\times 10^{-2}$ & 2.263 $\\times 10^{-2}$ & 4.419 $\\times 10^{-1}$ \\nl BLS1             & \\nodata & \\nodata                & 7.925 $\\times 10^{-1}$ & 2.515 $\\times 10^{-5}$ \\nl S1.5             & \\nodata & \\nodata                & \\nodata                & 5.194 $\\times 10^{-5}$ \\nl S2$_{\\rm total}$ & \\nodata & \\nodata                & \\nodata                & \\nodata                \\\\ \\enddata \\end{deluxetable}  In this section, we discuss why the observed values of $R_{\\rm O III}$ are higher in S1s than in S2s and how the high $R_{\\rm O III}$ comparing to the predicted values by simple photoionization models is achieved.  The first problem is the type dependence of the observed values of $R_{\\rm O III}$, shown in Figure 5. There are two possible interpretations to understand this dependence. One is that the [O {\\sc iii}]$\\lambda$5007 emission is stronger in S2s than S1s due to the intrinsically (i.e., not due to inclination effects) larger size of NLRs of S2s as proposed by Schmitt \\& Kinney (1996). Another is that the strong [O {\\sc iii}]$\\lambda$4363 emitting regions exist somewhere but obscured by something on the line of sight when we see S2s. Because the intrinsic AGN luminosity is similar among the various types of Seyferts in our sample as mentioned in Section 2.2, the former case predicts the stronger [O {\\sc iii}]$\\lambda$5007 luminosity in S2s than in S1s and similar [O {\\sc iii}]$\\lambda$4363 luminosity among the Seyfert types. On the other hand, the latter case predicts the similar [O {\\sc iii}]$\\lambda$5007 luminosity among the Seyfert types and the stronger [O {\\sc iii}]$\\lambda$4363 luminosity in S1s than in S2s. These two cases may be the extreme ones and the real situation might be intermediate between the two cases. However, it is interesting to investigate which case is close to the observed properties of emission-line spectra. Therefore we compare the observed emission-line luminosity of [O {\\sc iii}]$\\lambda$5007 and [O {\\sc iii}]$\\lambda$4363 among various types of Seyferts, which is shown in Figure 10. To quantify the statistical significance of the difference in the luminosity distribution among the Seyfert types, we apply the KS test where the null hypothesis is that the distribution of these emission-line luminosities come from the same underlying population. The KS test leads to the following results. (1) The distributions of the [O {\\sc iii}]$\\lambda$5007 luminosity are statistically indistinguishable among the BLS1s, the S1.5s, and the S2s though the [O {\\sc iii}]$\\lambda$5007 luminosity of the NLS1s are weaker than that of other Seyfert types. (2) The distributions of the [O {\\sc iii}]$\\lambda$4363 luminosity are statistically indistinguishable among the NLS1s, BLS1s, and the S1.5s though the [O {\\sc iii}]$\\lambda$4363 luminosity of the S2s are weaker than that of the BLS1s and the S1.5s. The KS probabilities are given in Table 6. These results are consistent with the latter case; i.e., the type dependence of the observed values of $R_{\\rm O III}$ is not due to the dilution by a more extended NLR in S2s but due to the enhancement of the [O {\\sc iii}]$\\lambda$4363 emission in S1s. The reason why the [O {\\sc iii}]$\\lambda$5007 luminosity is smaller in the NLS1s than in other types may be that NLS1s are the extreme objects on the ``eigenvector 1'' of Boroson \\& Green (1992). More explicitly, there is a relation that the weaker the [O {\\sc iii}]$\\lambda$5007 emission is, the narrower the FWHM of H$\\beta$ is (e.g., Boroson \\& Green 1992; Brandt \\& Boller 1998; Sulentic et al. 2000). Following this relation, NLS1s may tend to exhibit weak [O {\\sc iii}]$\\lambda$5007 emissions as shown in Figure 10.  The properties of emission-line width described in Section 3.4 also support the idea that the dependence of $R_{\\rm O III}$ on AGN types is not due to the dilution of $R_{\\rm O III}$ by extended low density gas in S2s but due to the obscuration of the strong [O {\\sc iii}]$\\lambda$4363 emitting region in S2s. The S1s exhibit the broad [O {\\sc iii}]$\\lambda$4363 comparing to [O {\\sc iii}]$\\lambda$5007 while the S2s do not so. This suggests the existence of the inner, strong [O {\\sc iii}]$\\lambda$4363 emitting region which is obscured in S2s. We, therefore, conclude that the dependence of $R_{\\rm O III}$ on AGN types is attributed to the obscuration effect.  Now we consider the following problems. {\\it Where is the [O {\\sc iii}]$\\lambda$4363 emitting region in AGNs?} And, {\\it how is the high $R_{\\rm O III}$ comparing to the predicted values by simple photoionization models achieved?} According to the current unified model of AGNs, it is natural to consider that the material obscuring the strong [O {\\sc iii}]$\\lambda$4363 emitting regions in S2s is dusty tori. If this is the case, the strong [O {\\sc iii}]$\\lambda$4363 emitting region may be either the inner surface of dusty tori described by Pier \\& Voit (1995) or the dense gas clouds near the central engine, which are obscured by the tori, though these two alternatives cannot be distinguished only by the statistical tests presented in this paper. This is consistent with the similarity of the location between the HINER and the strong [O {\\sc iii}]$\\lambda$4363 emitting region, which is suggested by the correlation between the intensity of the HINER emission and $R_{\\rm O III}$ (see Section 3.3) because large parts of the HINER emission is thought to arise from such dense clouds (Pier \\& Voit 1995; Murayama \\& Taniguchi 1998a, 1998b; Nagao et al. 2000c). The MIR properties described in Section 3.2 also support this geometrical consideration of the [O {\\sc iii}]$\\lambda$4363 emitting region; i.e., the correlations between the MIR colors and $R_{\\rm O III}$ mean the similarity between the visibility of relatively hot inner surface of dusty tori and that of the strong [O {\\sc iii}]$\\lambda$4363 emitting region.  However, some of S2s also exhibit large $R_{\\rm O III}$, which is also difficult to be explained by the simple one-zone photoionization model. They may be have dense gas clouds in NLR, which is not obscured by dusty tori because of the large distance from the nucleus, or the escaping [O {\\sc iii}]$\\lambda$4363 emission from the leaky parts of the tori.  \\subsection{Interpretation Using the Dual-Component Photoionization Model}  \\begin{figure*} \\epsscale{1.8} \\plotone{fig11.ps} \\caption{ The calculated emission-line ratio of [Fe {\\sc x}]$\\lambda$6374/[Fe {\\sc vii}]$\\lambda$6087 (solid line) and the column density (dashed line) of the truncated torus component are shown as functions of the ionization parameter of the torus component, $U_{\\rm DC}$. \\label{fig11}} \\end{figure*}  \\begin{figure*} \\epsscale{1.0} \\plotone{fig12.ps} \\caption{ The diagram of $R_{\\rm O III}$ versus [Fe {\\sc vii}]$\\lambda$4363/[O {\\sc iii}]$\\lambda$5007. The symbols are the same as in Figure 2. Our model calculations are superposed in the figure. The fraction of the contribution from torus components are also shown. The data points will move on the diagram as shown by the arrow if the extinction correlation of $A_V$ = 1.0 is applied. \\label{fig12}} \\end{figure*}  As described in Section 4.1, high-density gas clouds obscured by tori, which are located either at the inner surface of dusty tori or near the central engine where is obscured by the tori, are thought to emit a large fraction of the [O {\\sc iii}]$\\lambda$4363 emission. In order to investigate whether or not such an idea is consistent with photoionization scenarios quantitatively, we perform dual-component photoionization model calculations following the manner of Murayama \\& Taniguchi (1998b). This method takes account of such high-density gas clouds as a strong [O {\\sc iii}]$\\lambda$4363 emitter, in addition to the typical NLR component. Here we assume the second situation, i.e., the strong [O {\\sc iii}]$\\lambda$4363 emitting region is not physically associated by the tori. Therefore we do not consider effects of dust, such as depletion of heavy metals, through the following calculations.  Our calculation methods are as follows. We perform photoionization model calculations using the spectral synthesis code {\\it Cloudy} version 90.04 (Ferland 1996), which solves the equations of statistical and thermal equilibrium and produces a self-consistent model of the run of temperature as a function of depth into the nebula. Here we assume an uniform-density gas cloud with a plane-parallel geometry. The dense component (DC) is assumed to be truncated clouds; i.e., optically thin clouds for the ionizing photons, for the purpose of avoiding unusually strong [O {\\sc i}] emission (see Murayama \\& Taniguchi 1998b). The parameters for the calculations are (1) the hydrogen density of the cloud ($n_{\\rm DC}$ and $n_{\\rm NLR}$), (2) the ionization parameter ($U_{\\rm DC}$ and $U_{\\rm NLR}$), which is defined as the ratio of the ionizing photon density to the electron density, (3) the thickness of the torus component which is represented by the optical depth for ionizing photons, (4) the chemical compositions of the gas, (5) the shape of the input SED of ionizing photons, and (6) the fraction of DC to the NLR component.  Here we assume $n_{\\rm DC} = 10^7$ cm$^{-3}$. We perform several model runs covering 10$^1$ cm$^{-3} \\leq n_{\\rm NLR} \\leq 10^6$ cm$^{-3}$. The ionization parameter of the NLR component is assumed as $U_{\\rm NLR} = 10^{-2}$. The ionization parameter and the hydrogen column density of DC are determined using following two conditions; ([Fe {\\sc x}]$\\lambda$6374/[Fe {\\sc vii}]$\\lambda$6087)$_{\\rm DC}$ = 0.8 and ([Fe {\\sc vii}]$\\lambda$6087/[O {\\sc iii}]$\\lambda$5007)$_{\\rm DC}$ = 1.0. The former ratio is the typical value of Seyfert galaxies (Nagao et al. 2000c) and the latter condition is introduced by Murayama \\& Taniguchi (1998b) as a truncated dense gas cloud. As the result, $U_{\\rm DC} = 10^{-1.48}$ and $N_{\\rm DC} = 10^{20.76}$ cm$^{-2}$ are adopted \\footnote{ Because it is known that there is uncertainty in collision strengths for the [Fe {\\sc x}]$\\lambda$6374 emission, the derived values of $U_{\\rm DC}$ and $N_{\\rm DC}$ may also suffer such uncertainties. } (see Figure 10). The calculations are stopped when the gas temperature falls to 4000 K for the NLR component. We set the gas-phase elemental abundances to be solar ones taken from Grevesse \\& Anders (1989) with extensions by Grevesse \\& Noels (1993). We adopt the power-law continuum as the input spectrum, where the spectral index is assumed as $\\alpha$ = --1.5 (see Ferland \\& Netzer 1983) between 10 $\\mu$m and 50 keV for the form $f_{\\nu} \\propto \\nu^{\\alpha}$. The spectral index is set to $\\alpha$ = 2.5 at lower energy (i.e., $\\lambda \\geq 10 \\mu$m) and to $\\alpha$ = --2 at higher energy (i.e., $h\\nu \\geq$ 50 keV). The fraction of DC to the NLR component is treated as a free parameter in our calculations. The further details for this dual-component photoionization model are described in Murayama (1998) and Murayama \\& Taniguchi (1998b).  We present our results of model calculations and compare them with the observations in Figure 12, which is a diagram of $R_{\\rm O III}$ versus [Fe {\\sc vii}]$\\lambda$6087/[O {\\sc iii}]$\\lambda$5007. We find that the model grids are roughly consistent with the observations if we take the effects of the correction for the extinction into account. Though the dispersion of observation is larger than the model grids, this is thought to be attributed the fact that the parameters, such as $U_{\\rm DC}$, $U_{\\rm NLR}$, $n_{\\rm DC}$, and $N_{\\rm DC}$, are different from object to object. It is shown that the $R_{\\rm O III}$ of the S1s can be explained by introducing a 5\\% $\\sim$ 20\\% contribution from DC while the $R_{\\rm O III}$ of the S2s can be explained by introducing a 0\\% $\\sim$ 2\\% contribution from DC. These fractions are consistent with the results of Murayama \\& Taniguchi (1998b), who introduce a $\\sim 10 \\%$ contribution from the dense gas clouds to explain intensities of the HINER emission of S1s.  \\subsection{$R_{\\rm O III}$ in LINERs}  \\begin{figure*} \\epsscale{0.6} \\plotone{fig13.ps} \\caption{ The frequency distributions of $R_{\\rm O III}$ for the L1.9s and the L2s. \\label{fig13}} \\end{figure*}  \\begin{figure*} \\epsscale{1.7} \\plotone{fig14.ps} \\caption{ Same as Figure 9 but the sample of LINERs (NGC 3031, NGC 7213, and PKS 1718-649) are added. The crosses denote the LINERs. \\label{fig14}} \\end{figure*}  \\begin{deluxetable}{lll} \\tablenum{7} \\tablecaption{Observed values of $R_{\\rm OIII}$ for LINERs \\label{tbl-7}} \\tablehead{ \\colhead{Object Name} & \\colhead{$R_{\\rm OIII}$} & \\colhead{Reference\\tablenotemark{a}} } \\startdata \\cutinhead{LINER with a broad component} NGC 1052    & 0.0606 & HFS93 \\nl & 0.0353 & HFS97 \\nl NGC 1275    & 0.0923 & HFS93 \\nl NGC 3031    & 0.1884 & HFS93 \\nl & 0.0938 & HFS96 \\nl NGC 3226    & 0.0901 & HFS97 \\nl NGC 4278    & 0.1353 & HFS97 \\nl NGC 4395    & 0.0312 & HFS93 \\nl NGC 4579    & 0.0510 & HFS97 \\nl NGC 7213    & 0.2030 & FH84  \\nl \\cutinhead{LINER without a broad component} NGC 1167    & 0.0090 & HFS93 \\nl NGC 1961    & 0.0750 & HFS97 \\nl NGC 3504    & 0.0618 & HFS93 \\nl NGC 4102    & 0.0300 & HFS97 \\nl NGC 6500    & 0.0500 & HFS97 \\nl NGC 7714    & 0.0121 & HFS93 \\nl PKS 1718-649& 0.0960 & F85   \\\\ \\enddata \\tablenotetext{a}{ References for the data of $R_{\\rm OIII}$. Each abbreviation means as follows; F85: Filippenko (1985); FH84: Filippenko \\& Halpern (1984); HFS93: Ho, Filippenko, \\& Sargent (1993); HFS96: Ho, Filippenko, \\& Sargent (1996); and HFS97: Ho, Filippenko, \\& Sargent (1997a). } \\end{deluxetable}  In some low-ionization nuclear emission-line regions (LINERs), the observed values of $R_{\\rm O III}$ are far larger than that predicted by one-zone photoionization models. Because of this property, the dominant mechanism for the ionization in LINERs has been frequently regarded as shock ionization (e.g., Fosbury et al. 1978; Heckman 1980; Baldwin, Phillips, \\& Terlevich 1981). However, Filippenko (1985) pointed out that there is a correlation between the emission-line width and $n_{\\rm cr}$ of the emission line over the range 10$^3$ cm$^{-3} \\leq n_{\\rm cr} \\leq 10^7$ cm$^{-3}$ in a LINER PKS 1718-649. This suggests that LINERs may also possess high-density regions up to 10$^7$ cm$^{-3}$, which mean that the high $R_{\\rm O III}$ in LINERs may be explained by photoionization models.  Although it is not clear whether or not there is a dusty torus in all LINERs, there is several pieces of evidence that the unified models of AGNs can apply to some of LINERs; i.e., some LINERs exhibit broad components in their optical spectra (e.g., Ho et al. 1997b), in UV spectra (e.g., Barth et al. 1996), or only in polarized spectra (Barth, Filippenko, \\& Moran 1999a, 1999b). Therefore it is interesting to examine the properties of $R_{\\rm O III}$ of LINERs in the framework of our dual component model.  In order to investigate this issue, we compiled $R_{\\rm O III}$ of LINERs from the literature. Since the optical spectra of LINERs are often contaminated by stellar features from the host galaxy strongly, careful subtraction of such stellar features from observed spectra is needed to discuss the properties of faint emission lines such as [O {\\sc iii}]$\\lambda$4363. Our compiled sample consists of the objects that such careful subtraction was applied to; 8 LINERs with broad components (L1.9s) and 7 LINERs without broad components (L2s). Here it must be kept in mind that some of the L2s may not be AGNs; a part of LINERs may be shock-heated galaxy (e.g., Heckman 1980; Baldwin et al. 1981; Heckman 1986; Gonz\\'arez Delgado \\& P\\'erez 1996) and others may be the objects ionized by hot stellar component (e.g., Filippenko \\& Terlevich 1992; Binette et al. 1994; Alonso-Herrero et al. 2000; Taniguchi, Shioya, \\& Murayama 2000). It is noted that the compiled samples may be biased in favor of the higher $R_{\\rm O III}$ objects because it is often difficult to detect weak [O {\\sc iii}]$\\lambda$4363 emission owing to the relatively strong stellar feature. The objects we compiled are given in Table 7, and the histograms of $R_{\\rm O III}$ for the L1.9s and the L2s are shown in Figure 13. The means and the deviations of $R_{\\rm O III}$ are 0.099 $\\pm$ 0.054 for the L1.9s and 0.048 $\\pm$ 0.030 for the L2s. There are a tendency that the values of the L1.9s are higher than those of L2s though these are statistically indistinguishable ($P_{\\rm KS}$ = 0.252).  In the sample, the emission-line width of [O {\\sc iii}]$\\lambda$4363 has been measured for three LINERs; NGC 3031 (Ho, Filippenko, \\& Sargent 1996), NGC 7213 (Filippenko \\& Halpern 1984), and PKS 1718-649 (Filippenko 1985). These data follow our dual component model; i.e., higher $R_{\\rm O III}$ objects show larger values of FWHM([O {\\sc iii}]$\\lambda$4363) and FWHM([O {\\sc iii}]$\\lambda$4363)/FWHM([O {\\sc iii}]$\\lambda$5007 (Figure 14).  These results seem to suggest that the [O {\\sc iii}]$\\lambda$4363 emitting regions are located at inner than the [O {\\sc iii}]$\\lambda$5007 emitting regions and that the [O {\\sc iii}]$\\lambda$4363 emission has an anisotropic property for LINERs, too. Further observations are needed to investigate this issue in detail."
    },
    "0011/astro-ph0011333_arXiv.txt": {
        "abstract": "I review the observational data most relevant for large scale structure.  These data determine the system of cosmological parameters: the Hubble parameter, densities of various populations of the Universe, parameters characterizing the power spectrum of matter, including the biasing parameter of galaxies relative to matter. Recent data suggest that the overall matter/energy density is approximately equal to the critical density, and most ($0.6 - 0.7$) of the density is in the form of cosmological term or ``dark (vacuum) energy''.  The density of the matter is $0.3 - 0.4$ (including hot and cold dark matter and luminous matter), the upper limit of the density of the hot dark matter is 0.05, all in units of the critical cosmological density. ",
        "introduction": "Recent results from the supernova cosmology project \\cite{perl98}, \\cite{riess98} and measurements of the angular spectrum of the cosmic microwave background (CMB) radiation \\cite{boomerang}, \\cite{maxima} have triggered a number of efforts to determine a concordant system of cosmological parameters.  In this talk I shall use recent observational data to discuss values of main cosmological parameters. In addition to data on the CMB angular spectrum and supernova cosmology project I shall use data based on the large-scale distribution of galaxies and clusters of galaxies -- the power spectrum, the cluster mass function etc.  In this analysis I use the following assumptions: 1) the main constituents of the Universe are baryonic matter, cold dark matter (CDM) with some mixture of hot dark matter (HDM), and the dark (vacuum) energy; 2) power spectra of galaxies and CMB radiation are determined by the initial post-inflational power spectrum and by physical processes during the radiation--dominated era. These processes depend on cosmological parameters (properties of various components of the matter), and on geometrical properties of the Universe.  In this analysis I try to find the possible range of cosmological parameters and to show how these are affected by the various types of data. ",
        "conclusions": "BOOMERANG and MAXIMA I data have been used in a number of studies to determine cosmological parameters \\cite{bz00}, \\cite{cen00}, \\cite{boomerang}, \\cite{maxima}, \\cite{kinney00}, \\cite{tz00}, \\cite{wsp00}.  In addition to CMB data various other observational data have been used.  In general, the agreement between various determinations is good; however, some parameters differ.  For instance, \\cite{tz00} interpreted new CMB data in terms of a baryon fraction higher than expected from the nucleosynthesis constraint, $h^2 \\Omega_b = 0.03$, and a relatively high matter density, $h^2 \\Omega_m = 0.33$.  On the other hand, velocity data suggest a relatively high amplitude of the power spectrum, $\\sigma_8 \\Omega_m^{0.6} = 0.54$, which in combination with distant supernova data yields $\\Omega_m = 0.28 \\pm 0.10$ and $\\sigma_8 = 1.17 \\pm 0.2$ \\cite{bz00}.  Our analysis has shown that a high value of the density of matter, $\\Omega_m >0.4$, and high amplitude of the matter power spectrum, $\\sigma_8 >1$, are difficult to explain in terms of the supernova and cluster abundance data, and the observed amplitude of the galaxy power spectrum with reasonable bias limits.  This conflict can be avoided using a tilted initial power spectrum, and a MDM model with a moderate fraction of HDM, as discussed above. The best models suggested so far have $0.3 \\leq \\Omega_m \\leq 0.4$, $0.90 \\leq n \\leq 0.95$, $0.60 \\leq h \\leq 0.70$, $\\Omega_n \\leq 0.05$.  Matter densities are constrained to $\\ge 0.3$ by cluster abundances, and to $\\le 0.4$ by all existing matter density estimates. This upper limit of the matter density, in combination with the cluster abundance and amplitude of the power spectrum, yields an upper limit to the density of the hot dark matter. We can consider this range of cosmological parameters as compatible with all constraints.  This set of cosmological parameters is surprisingly close to the set suggested by Ostriker \\& Steinhardt \\cite{os95}.   Now it is supported by much more accurate observational data.   A considerably lower value of matter density, $\\Omega_m = 0.16$, was suggested by Bahcall \\etal \\cite{b00} from the observed value of $M/L$ for galaxies and clusters of galaxies of various richness. Upper right panel of Fig.~1 shows this constraint for various fractions of matter in voids and respective bias parameter values. The reason for the deviation of this matter density determination from the rest is not clear, and we have not used it in the present analysis."
    },
    "0011/astro-ph0011043_arXiv.txt": {
        "abstract": "The formation of [O\\,{\\sc i}], O\\,{\\sc i} and OH lines in metal-poor stars has been studied by means of 3D hydrodynamical model atmospheres. For O\\,{\\sc i} detailed 3D non-LTE calculations have been performed. While the influence of 3D model atmospheres is minor for [O\\,{\\sc i}] and O\\,{\\sc i} lines, the very low temperatures encountered at low metallicities have a drastic impact on OH. As a result, the derived O abundances are found to be systematically overestimated in 1D analyses, casting doubts on the recent claims for a monotonic increase in [O/Fe] towards lower metallicities. \\footnote{The full version of the present talk with additional figures can be found at http://www.astro.uu.se/$\\sim$martin/talks/Manchester00} ",
        "introduction": "Besides being relatively common, oxygen has attracted a great deal of attention due to its special role in galactic and stellar evolution. The ratio [O/Fe] at different cosmic epochs give important insight to the formation and evolution of the Galaxy, as well as constraining the physics of supernovae. Oxygen can also significantly influence stellar evolution and is crucial for the production of the light elements Li, Be and B through cosmic ray spallation. Following the first indication of an oxygen over-abundance relative to iron in metal-poor stars, a great number of studies have been devoted to quantifying this enhancement but with little concordance. The forbidden [O\\,{\\sc i}] line at 630.0\\,nm in metal-poor giants suggest a nearly flat plateau at [O/Fe]\\,$\\sim 0.5$ for [Fe/H]\\,$\\la -1.0$ while the O\\,{\\sc i} triplet at 777\\,nm in dwarfs tend to imply systematically higher values, often with a monotonic increase towards lower metallicities. Recently, Israelian et al. (1998) and Boesgaard et al. (1999) have analysed the UV OH lines down to [Fe/H]~$\\simeq -3.3$ and have found a linear trend in [O/Fe] vs [Fe/H] with a slope of about -0.4, in stark conflict with the [O\\,{\\sc i}] results. All oxygen criteria have their pros and cons, which influence the conclusions. Of particular interest is the possible systematic errors introduced by the inherent assumptions behind the 1D model atmospheres employed in the analyses of the various O diagnostics. Significant progress in this respect has recently been accomplished with the construction of 3D stellar convection simulations (e.g. Stein \\& Nordlund 1998; Asplund et al. 1999, 2000a). The aim of the present investigation is to study the impact of such improved models on oxygen line formation. ",
        "conclusions": "The new generation of 3D hydrodynamical model atmospheres has a dramatic impact on in particular the OH line formation in metal-poor stars. Caution must, however, be exercised in interpreting these 3D findings in terms of galactic cosmic evolution until detailed 3D non-LTE calculations for the OH molecular formation and line formation have been performed."
    },
    "0011/astro-ph0011569_arXiv.txt": {
        "abstract": "A radio-based search for strong gravitational lensing, with image separations in the range 160--300 milliarcsec (mas), has yielded a null result for a sample of 1665 sources (Augusto, Wilkinson \\& Browne 1998) whose mean redshift is estimated to be $\\sim 1.3$. The lensing rate for this previously-unexplored separation range, $<1:555$ at the 95\\% confidence level, is less than on arcsecond-scales---as expected from models of lensing galaxy populations. Lensing on 160--300 mas scales is expected to arise predominantly from spiral galaxies at a rate dependent on the disk-halo mass ratio and the evolving number density of the population with redshift. While the present sample is too small for there to be a high probability of finding spiral galaxy lenses, our work is a pilot survey for a  much larger search based on the full CLASS database which would provide useful information on galactic structure at $z \\sim 0.5$. We examine other possible lens populations relevant to our present search, in particular dwarf galaxies and supermassive black holes in galactic nuclei, and conclude that none of them are likely to be detected. Our null result enables us formally to rule out a cosmologically significant population of uniformly-distributed  compact objects: $\\Omega_{\\rm CO}<0.1$ (95\\% confidence) in the mass range $10^{9.5}$--$10^{10.9}$~M$_{\\odot}$. ",
        "introduction": "Systematic radio-based surveys have proved to be a very successful way of finding gravitational lens systems with image separations in the range 0.3--6 arcsec; the lensing masses are the central regions of normal galaxies.  The Jodrell-VLA Astrometric Survey (JVAS)/Cosmic Lens All-Sky Survey (CLASS) surveyed over 12000 sources, using the VLA, MERLIN and the VLBA as successive high resolution filters. These have, so far, found 19 such lens systems (Browne et al.\\ 2000 and unpublished). In the course of a general investigation of the properties of flat-spectrum radio sources, drawn principally from the JVAS, Augusto, Wilkinson \\& Browne (1998; paper I hereafter) carried out a search for multiple imaging with smaller separations in the range $\\sim$90--300 mas. Using essentially the same VLA, MERLIN and VLBA filtering process employed in the main JVAS/CLASS lens surveys, no cases of multiple imaging were found in a sample of 1665 sources. In this paper we examine the implications of this null result.  The classic statistical study of gravitational lensing probabilities was carried out by Turner, Ostriker \\& Gott (1984; TOG84 hereafter) who calculated the distribution of the angular separations of multiple images expected from normal galaxies. TOG84 predicted that $\\sim 88\\%$ of lenses associated with normal galaxies should occur in the range 0.3--6 arcsec. All of the known cases of strong lensing by individual galaxies fall in this angular size range, with a mean separation consistent with TOG84's expectation of $\\sim 1.5$ arcsec (Browne et al.\\ 2000) and a smallest separation of 0.335 arcsec \\cite{Patetal93}. It is important, however, to subject the TOG84 predictions to further observational scrutiny in order to see whether there are other contributions to the lens population. Our results are for a previously unexplored separation range and hence provide a test for lower-mass galaxies (or components of galaxies) than were considered by TOG84. This is particularly relevant since recent calculations of the lensing contribution of spiral galaxies suggest that TOG84 may have underestimated their contribution to the lensing optical depth by more than a factor two (see section 2.1).  Our results also enable us to place observational constraints on `non-standard' populations of potential lensing masses. Any mass concentration with size $\\la0.02 \\sqrt{M/M_{\\odot}}$ pc is `compact' as regards gravitational lensing \\cite{PreGun73} and unseen dark Compact Objects (CO) could produce lensing on scales of interest for this paper. If present in large enough numbers, such CO could contribute to the solution of the dark matter problem.  In the extreme case of a universe filled with a critical density of CO of a particular mass ($\\Omega_{\\rm CO}=1$), the odds are even that any object at $z=2$ will be multiply imaged into two images with a brightness ratio $<10:1$ \\cite{PreGun73}. The lensing probability scales directly with the number density of CO and hence the statistics of gravitational lensing in surveys of distant sources provide direct constraints on $\\Omega_{\\rm CO}$.  Throughout this paper we take the value of $H_0$ to be 65 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Among a sample of 1665 flat-spectrum radio sources, with an assumed mean redshift of $\\sim 1.3$, no multiple images were found with separations in the range 160--300 mas (paper I and Augusto, unpublished). The lensing rate for compact masses in the range $\\sim 10^{9.5}$--$10^{10.9}$~M$_{\\odot}$ is therefore less than $1:555$ at 95\\% confidence. This result does not contradict TOG84 predictions based on the properties of normal E/S galaxies.  No unexpected population of small galaxies has thus been detected and it seems that, despite their large numbers, $\\sim10^9$ M$_{\\odot}$ dwarf galaxies are not contributing significantly to multiple imaging on 160-300 mas. This is consistent with the current knowledge about their numbers and mass distributions.  Recent work on the contribution of $z \\sim 0.5$ spirals to the overall lensing statistics gives contradictory results: there could be significant (greater than a factor of two) enhancements of the total spiral fractional contribution.  Part of the motivation for this new work on the spiral population has been to anticipate the capabilities of ALMA and NGST at $\\sim$0.1 arcsec resolution in the sub-mm and infra-red bands.  Radio-based surveys are similarly well-suited to an unbiased search for spiral galaxy lensing since they are also not affected by dust obscuration in edge-on disk systems. It is now possible to mount a significant test of this prediction, since by using the same methodology as adopted for Paper~I we can look for compact lenses in the now-completed CLASS surveys.  Since the combined JVAS/CLASS surveys have, so far, found 19 lenses with separations 0.3--6 arcsec in a total sample of $\\sim12000$ sources, a search in the range 0.1-0.3 arcsec using the same parent sample should find some small-separation lenses (likely to be spirals). The results would place significant contraints on the uncertain properties of spiral galaxies at the redshifts appropriate to multiple imaging ($z\\sim 0.5$).        Analysing the simpler point mass lens constraints, we conclude that the cosmological density of compact objects  (which include black holes and $\\sim$1--5 kpc compact galaxies) in the mass range $\\sim3\\times10^{9}$ to $\\sim8\\times10^{10}$ M$_{\\odot}$, has to be $\\Omega_{CO}^{9.5-10.9}<0.10$ (95\\% confidence level). From our new result, it seems that $10^{9.5}$--$10^{10.9}$ M$_{\\odot}$ SBH in this range cannot contribute significantly to solving the dark matter problem."
    },
    "0011/astro-ph0011275_arXiv.txt": {
        "abstract": "We study the globular cluster (GC) system of the dust-lane elliptical galaxy NGC\\,6702, using $B$, $V$ and  $I$-band photometric observations carried out at the Keck telescope. This galaxy has a spectroscopic age of $\\approx2$\\,Gyrs suggesting recent star-formation. We find strong evidence for a bimodal GC  colour distribution, with the blue peak having a colour similar to that of the Galactic halo GCs. Assuming that the blue GCs are indeed old and metal--poor, we estimate an age of 2--5\\,Gyrs and supersolar metallicity for the red GC subpopulation. Despite the large uncertainties, this is in reasonable agreement with the spectroscopic galaxy age. Additionally, we estimate a specific frequency of $S_{N}=2.3\\pm1.1$ for NGC\\,6702. We predict that passive evolution of NGC\\,6702 will further increase its specific frequency to  $S_{N}\\approx2.7$ within 10\\,Gyrs, in closer agreement to that of  typical present-day ellipticals.  We also discuss evidence that the merger/accretion event that took place a few Gyrs ago involved a high gas fraction. ",
        "introduction": "The `merger' hypothesis, postulating that elliptical galaxies form by spiral galaxy mergers, was first proposed by Toomre \\& Toomre (1972). This scenario has gained new impetus over the last few years by both numerical simulations (e.g. Barnes 1992; Hernquist 1992; Mihos \\& Hernquist 1996) and observational studies, suggesting that the properties of merger remnants are consistent with those of ellipticals (e.g. Toomre 1977;  Schweizer et al. 1990;  Casoli et al. 1991; Schweizer \\& Seitzer 1992; Hibbard \\& van Gorkom 1996; Georgakakis, Forbes \\& Norris 2000).  However, one of the long standing drawbacks of the `merger' hypothesis is the excess globular clusters (GCs) per unit starlight around ellipticals compared to spirals (van den Bergh 1984). To overcome this problem Schweizer (1987) and Ashman \\& Zepf (1992) suggested that the starburst activity usually experienced during gas-rich mergers will form new GCs and thus increase their number compared to their progenitors.   Indeed, there is accumulating observational evidence supporting this scenario in on-going mergers (e.g. Lutz 1991; Whitmore et al. 1993; Whitmore \\& Schweizer 1995; Schweizer 1996; Miller et al. 1997; Zepf et al. 1999; Forbes \\& Hau 2000). However, the situation in ellitpical galaxies (the eventual product of mergers) is less clear. Although many ellipticals have been observed to have bimodal colour distributions, when examined in detail their GC properties do not match those from the merger picture (Forbes et al. 1997). Outstanding questions include whether the newly formed GCs survive and whether the final GC system resembles that of typical (old) ellipticals. Probing the GC properties in intermediate age merger-remnants (1--5\\,Gyrs) may be of key importance for addressing these issues. Indeed, such young ellipticals may provide the link between on-going mergers and relaxed ellipticals.  For example, a sample of merger--remnants spanning a range of ages will allow study of the chronological development of the system's specific frequency (measuring the number of globular clusters per unit starlight) in comparison with that of evolved ellipticals. Young ellipticals are also expected to have two distinct GC populations with different colour and magnitude distributions. By exploiting these differences one can attempt to age-date the newly formed GCs, which in turn can be compared to the time since the gaseous merger. Indeed, studies of the  dynamically young proto-ellipticals NGC\\,1700 and NGC\\,3610  have shown the presence of bimodality in their colour distributions (Whitmore et al. 1997; Brown et al. 2000). Also the age of the new GCs, although model dependent, is found to be similar to estimates for the galaxy age based on spectroscopic or dynamical arguments.   In this paper we study the globular cluster system of the elliptical galaxy NGC\\,6702. Lauer (1985) found a dust lane in the centre of NGC\\,6702 indicating a recent gaseous merger event. This galaxy is also relatively luminous  at far-infrared wavelengths compared to other elliptical galaxies, with  $L_{FIR}=1.3\\times10^{10}\\,\\mathrm{L_{\\odot}}$  (Wiklind et al. 1995).  Terlevich \\& Forbes (2000) have estimated a spectroscopic age and metallicity of about 2\\,Gyrs and $\\mathrm{[Fe/H]=+0.5}$  respectively for NGC\\,6702 suggesting that it is a young elliptical. The main observational properties of NGC\\,6702 are given in Table \\ref{tab_0}.  Additionally, according to the NED database the nearest galaxy projected on the sky that is either brighter or no more than 2 magnitudes fainter, than NGC\\,6702 is  NGC\\,6703 at a projected distance of 10\\,arcmin. However, as NGC\\,6702 has a recession velocity of $4727\\,\\mathrm{km\\,s^{-1}}$ and NGC\\,6703 has $V=2365\\mathrm{\\,km\\,s^{-1}}$ they are actually separated by some 30\\,Mpc  assuming no peculiar motions. The nearest galaxy with a similar recession velocity is NGC\\,6711 ($V=4671\\mathrm{\\,km\\,s^{-1}}$) with a separation of $\\sim120$\\,arcmin or over 2\\,Mpc. Other galaxies listed by NED are more than 2 magnitudes fainter than NGC\\,6702. Thus NGC\\,6702 is an example of an isolated elliptical galaxy with no comparable galaxies within 2\\,Mpc.  Such galaxies are fairly rare and hence their evolutionary history is of some interest. In particular, ellipticals in low density environments are interesting within the hierarchical merger framework as they are expected to have assembled much more recently than their cluster counterparts (Kauffmann \\& Charlot 1998). Any merger would be expected to have occurred in the more recent past and involved progenitors that have largely processed gas into stars.  In section \\ref{sec_obs} the observations, data reduction and  the isophotal properties of NGC\\,6702 are presented. Section \\ref{sec_sel} describes the candidate GC selection, while section \\ref{sec_res} presents our results which are discussed in section \\ref{sec_disc}. Finally,  section \\ref{sec_conc} summarises our conclusions. Throughout this  paper we adopt $H_0=75\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$.  \\begin{table*} \\footnotesize \\begin{center} \\begin{tabular}{cccccccccc} \\hline Name  &  RA (J2000) & DEC (J2000) & l & b & type & heliocentric& distance & $B_T^0$ & $V_T^0$ \\\\ &             &             &   &   & (RC3)& velocity ($\\mathrm{\\,km\\,s^{-1}}$)  &  modulus &  (mag)    &  (mag)\\\\ NGC\\,6702 &  $18\\mathrm{^h}$ $46\\mathrm{^m}$ $57.6\\mathrm{^s}$ &  $+45\\mathrm{^\\circ}$ $42\\mathrm{^\\prime}$ $20\\mathrm{^{\\prime\\prime}}$ &  $75.03\\mathrm{^\\circ}$ &  $19.79\\mathrm{^\\circ}$ &  E\\,3 &  4727 &  34.00 &  13.04 &  12.15 \\\\ \\hline \\end{tabular} \\end{center} \\caption{Observational properties of NGC\\,6702 obtained from NED. The distance modulus is estimated assuming $H_0=75\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$.}\\label{tab_0} \\normalsize \\end{table*} ",
        "conclusions": "\\label{sec_conc}  We have carried out $B$, $V$ and  $I$-band photometric observations of the elliptical galaxy NGC\\,6702 to explore its GC system. The GC candidates are selected on the basis of their colour, magnitude and morphology. We find strong evidence that the $B-I$ colour distribution of GC candidates is bimodal with the blue peak having a colour similar to that of the  old metal--poor Galactic GCs. Assuming that the blue GC subpopulation is indeed old and metal--poor, we estimate that the red GCs have an age of 2--5\\,Gyrs and super-solar metallicity. Despite the large uncertainties, this is in reasonable agreement with the spectroscopically estimated age and metallcity of about 2\\,Gyrs and 0.5 respectively for the galaxy. Thus the red GCs in NGC\\,6702 are consistent with having formed during the most recent starburst, likely to have been triggered by a merger or an accretion event, a few Gyrs earlier. Additionally, we find a  specific frequency for the NGC\\,6702 GC system of $S_N=2.3\\pm1.1$. We predict that passive evolution  of the NGC\\,6702 stellar population formed during the most recent star-formation event will increase the NGC\\,6702 specific frequency to $\\approx2.7$ within 10\\,Gyrs, in reasonable agreement with the mean for present-day ellipticals, excluding the largest cluster ellipticals. Finally, using the NGC\\,6702 GC properties and making reasonable assumptions about the GC formation efficiency, we estimate the mass fraction of newly formed stars during the most recent star-formation event. We find reasonable agreement with the mass fraction estimated using stellar population synthesis models to interpret the line indices of the underlying galaxy starlight."
    },
    "0011/astro-ph0011039_arXiv.txt": {
        "abstract": "The application of Voronoi and Delaunay tessellation based methods for reconstructing continuous fields from discretely sampled data sets is discussed. The succesfull operation as ``multidimensional interpolation'' method is corroborated through their ability to reproduce even intricate statistical aspects of the analytical predictions of perturbation theory of cosmic velocity field evolution. The newly developed and fully self-adaptive technique for density field estimation by means of Delaunay tessellations, basically exploiting their ``minimum triangulation'' quality, is shown to succesfully reproduce the morphology of the foamlike structure in N-body simulations of cosmic structure formation. The full hierarchy of structure is implicitly and directly reproduced at every spatial resolution scale present in the particle distribution, at the same time automatically and sharply rendering its characteristic anisotropic filamentary and wall-like features. ",
        "introduction": "Astronomical observations, physical experiments as well as computer simulations often involve discrete data sets supposed to represent a fair sample of an underlying smooth and continuous field. Reconstructing the underlying fields from a set of irregularly sampled data is therefore a recurring key issue in operations on astronomical data sets.  Within the context of the reconstruction issue, we may distinguish two basically distinct situations. One is that of a specific continuous field whose values have been measured at a set of discrete locations. A typical example, of a cosmological nature, concerns the sampling of the global cosmic matter flow involved in the build up of structure in the Universe. The measured peculiar velocities of galaxies are supposed to be a fair reflection of the underlying cosmic flow. The reconstruction problem may then be described as  an issue of ``Multidimensional Interpolation''. Evidently, a myriad of astronomical studies involve related  such issues. A second class involves the issue of estimating the underlying continuous intensity or density field from a point process supposedly representing a fair sampling of this field. The fact that the sampling point process itself is a reflection of the underlying field forms an extra complication in the case of these ``Optimal Density Estimation'' problems. \\begin{figure*} \\includegraphics*[width=\\hsize,keepaspectratio]{mpa2000weyrf1.eps} \\caption{\\small The Delaunay triangulation (left frame) and Voronoi tessellation (right frame) of a distribution of 25 nuclei (stars) in a square (central panel). Periodic boundary conditions are assumed.} \\end{figure*} Conventional reconstruction methods are usually plagued by one or more artefacts. Firstly, they often involve estimates at a finite number of locations, usually confined to a grid. Optimal field estimators should provide a prescription for the value of that field throughout the whole sampling volume. Conventional schemes usually restrict themselves to estimates at a finite number of positions. For most practical purposes, a further disadvantage of almost all conventional methods is their insensitivity and inflexibility to the sampling point process. This leads to a far from optimal performance in both high density and low density regions, which often is dealt with by rather artificial and ad hoc means.  In particular in situations of highly non-uniform distributions conventional methods tend to obscure various interesting and relevant aspects present in the data. An illuminating example is the case of the galaxy distribution and the cosmic matter density field. Redshift surveys as well as computer simulations of structure formation reveal salient anisotropic features like filaments and walls, extended along one or two directions while compact in the other(s). In addition, the density fields display hierarchical structure of varying contrasts over a large range of scales. Ideally sampled by the data points, appropriate field reconstructions should be set solely and automatically by the sampling point distribution itself. The commonly used methods, involving artificial filtering through grid size or other smoothing kernels (e.g. Gaussian filter) mostly fail to achieve an optimal result and are neither able to reproduce the tenuous anisotropic structures seen in the point process, nor a faithful rendering of the full structural hierarchy. A final, usually concealed, yet fundamental aspect of field reconstruction turns out to be of key significance for the definition of the presented tessellation methods. Most well-known methods implicitly yield mass-weighted averages. This renders the comparison with volume-weighted analytical quantities far from trivial. Tessellation methods represent an elegant solution to this point. ",
        "conclusions": ""
    },
    "0011/astro-ph0011513_arXiv.txt": {
        "abstract": "We show that the difference between the Ar and Si relative abundance ratio derived from {\\it FUSE} absorption spectra and from  the \\ion{H}{ii} regions of I~Zw~18 is a consequence of the microturbulent analysis applied to the absorption spectra. {\\it FUSE} observations were performed with a large entrance aperture which fully covered the galaxy. This means that the observed profiles are averaged over the full body of I~Zw~18, implying that large-scale velocity fields influence the absorption-line profiles. Taking this into account, we show that the absorption spectra are consistent with the same metal abundances as those derived from the \\ion{H}{ii} regions. It follows that no significant ionization correction as suggested by Izotov and collaborators to describe metal contents in damped Ly$\\alpha$ systems (DLA) is required to model abundances in the neutral gas of I~Zw~18 (a local DLA system). Using a mesoturbulent approach and applying the generalized radiative transfer equation to the \\ion{Ar}{i}\\,$\\lambda1048$ and \\ion{Si}{ii}\\,$\\lambda1020$ lines observed by Vidal-Madjar et al., we found that the profiles may be reproduced with log (Ar/Si) $\\simeq - 0.8$ and $N(\\ion{Si}{ii}) \\simeq 4\\times10^{15}$ cm$^{-2}$. ",
        "introduction": "The blue compact galaxy (BCG) I~Zw~18 (Mrk 116) has been an intensively studied object for the last three decades since the first spectroscopic observations by Zwicky (1966). I~Zw~18 shows an intense and recent burst of star formation which makes this galaxy an attractive target for studies of star formation history. Recent observations revealed, for example, two stellar populations in I~Zw~18: 10--20 Myr red supergiants and 0.1--5 Gyr asymptotic giant branch stars (\\\"Ostlin 2000).  Amongst BCGs, I~Zw~18 shows the lowest oxygen abundance~: the northwest (NW) and the southeast (SE) bright \\ion{H}{ii} regions yield log (O/H) $= -4.83 \\pm 0.03$ and $-4.82 \\pm 0.03$, respectively (Izotov et al. 1999). Relative to solar values, one finds [O/H]\\footnote{Using the customary definition [X/H] = $\\log [N({\\rm X})/N({\\rm H})] - \\log [N({\\rm X})/N({\\rm H})]_\\odot$, and solar abundances from Grevesse et al. (1996) except for Ar for which the weighted average value $-5.48 \\pm 0.04$ from Sofia \\& Jenkins (1998) is adopted.} $ = -1.70 \\pm 0.08$, i.e. $Z/Z_\\odot \\simeq 1/50$. For argon abundances, Izotov et al. obtained log (Ar/H) $= -6.90 \\pm 0.05$ (NW) and a slightly lower ratio for the SE region, log (Ar/H) = $-7.23 \\pm 0.05$. Silicon abundances measured by Izotov \\& Thuan (1999) give log (Si/H) $= -6.29 \\pm 0.22$ (NW) and log (Si/H) $= -6.30 \\pm 0.22$ (SE). Thus, these values imply for the NW emission patch [Ar/H] = $-1.42 \\pm 0.06$ and [Si/H] = $-1.85 \\pm 0.22$.  \\begin{figure*} \\vspace{-1.5cm} \\hspace{0.0cm}\\psfig{figure=figur1.eps,height=9.0cm,width=18.0cm} \\vspace{-2.0cm} \\caption[]{({\\bf a}) -- Observed normalized intensities (dots with error bars) vs radial velocities from {\\it FUSE} spectra of \\ion{Ar}{i} (top panel) and \\ion{Si}{ii} (bottom panel) obtained by Vidal-Madjar et al. (2000). Smooth grey lines are the synthetic spectra (convolved with the spectrograph function) of the two-component microturbulent model calculated from the $\\chi^2$ minimization as described in the text. Smooth dotted lines show the corresponding unconvolved spectra. Components at $\\Delta v \\simeq 49$ km~s$^{-1}$ are the Galactic H$_2$ lines [\\,L(4-0)P(1) in top panel and L(7-0)P(4) in bottom panel\\,] from the high-velocity cloud at $-160$ km~s$^{-1}$. ({\\bf b}) -- Same as {\\bf a} but for a one-component microturbulent model. ({\\bf c}) -- Confidence regions in the $\\chi^2_{\\rm min} - b$ plane calculated from the simultaneous fit of the \\ion{Ar}{i} and \\ion{Si}{ii} lines shown in panels {\\bf b}. The parabola vertex corresponds to the best value of $b = 17.4$ km~s$^{-1}$ } \\end{figure*}  The neutral gas properties in I~Zw~18 have been probed in both the radio (see van Zee et al. 1998 and references cited therein) and the UV range (Kunth et al. 1994; Pettini \\& Lipman 1995; Vidal-Madjar et al. 2000). High velocity and high spatial resolution radio observations have shown that the overall kinematics of the \\ion{H}{i} gas associated with I~Zw~18 is very complex and the neutral gas velocity dispersion $\\sigma$ equals $13 - 14$ km~s$^{-1}$, or $b \\equiv \\sqrt{2}\\sigma = 18 - 20$ km~s$^{-1}$  (van Zee et al. 1998). The \\ion{H}{i} column density in front of I~Zw~18 deduced from the Ly$\\alpha$ absorption profile by Kunth et al., $N(\\ion{H}{i}) = (3.5 \\pm 0.5)\\times10^{21}$ cm$^{-2}$, is comparable to the peak \\ion{H}{i} surface density found by van Zee et al., $N(\\ion{H}{i}) \\simeq 3.0\\times10^{21}$ cm$^{-2}$. This means that the neutral gas in I~Zw~18 can be considered as a local damped Ly$\\alpha$ system (DLA) which is similar to high redshift DLAs observed in the light of background quasars.  First measurement of the O/H abundance in the \\ion{H}{i} region have indicated a possible discrepancy between the metal content in the neutral gas and in the \\ion{H}{ii} regions (Kunth et al. 1994). This, however,  was not confirmed in later studies by Pettini \\& Lipman (1995) and  by van Zee et al. (1998) who have shown that both the neutral and ionized gas in I~Zw~18 may have the same oxygen abundance.  New observations carried out with the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) by Vidal-Madjar et al. (2000) produce a similar puzzle~: the column density ratio deduced from the \\ion{Ar}{i}\\,$\\lambda1048$ and \\ion{Si}{ii}\\,$\\lambda1020$ lines in the neutral gas, log~(Ar/Si) $= -1.32$, differs significantly from that observed in the \\ion{H}{ii} regions, log~(Ar/Si)$_{\\rm NW} = -0.61 \\pm 0.22$ and log~(Ar/Si)$_{\\rm SE} = -0.74 \\pm 0.22$ (Izotov et al. 2000). Izotov et al. suggested that ionization in DLAs may affect abundance ratios. To interpret the observations, they suggested a model consisting of two regions with {\\it substantially different} metal contents, i.e. implicitly recalling an idea of the existence of two media in BCGs (a pristine gas, unprocessed since the big bang and a gas polluted by nucleosynthetic products -- Kunth \\& Sargent 1986). In general, the discovery of such a primordial gas would be of great importance for cosmology, as noted by Kunth et al. (1994).  In this Letter, we report on the study of the \\ion{Ar}{i}\\,$\\lambda1048$ and \\ion{Si}{ii}\\,$\\lambda1020$ profiles of I~Zw~18 published by Vidal-Madjar et al. (2000). We show that these lines can be modeled under the assumption that the metal content in the neutral gas is the same as in the \\ion{H}{ii} regions without refering to an ionization correction. ",
        "conclusions": "We have shown that metal abundances derived from the absorption profiles of \\ion{Ar}{i}\\,$\\lambda1048$ and \\ion{Si}{ii}\\,$\\lambda1020$, observed in the spectrum of the damped Ly$\\alpha$ system I~Zw~18 by Vidal-Madjar et al. (2000), become fully consistent with those derived from emission-line spectra of the \\ion{H}{ii} regions by Izotov \\& Thuan (1999) and by Izotov et al. (1999), if correlations in the large-scale velocity field are accounted for. The adequacy of the obtained mesoturbulent solutions is supported by the fact that the derived velocity dispersion of the neutral gas, $\\sigma \\simeq 13$ km~s$^{-1}$, is in an excellent agreement with the \\ion{H}{i} 21~cm observations, $\\sigma = 12 - 14$ km~s$^{-1}$ (van Zee et al. 1998). This implies that the gas in I~Zw~18 is efficiently mixed.  To estimate more accurately metal abundances in the \\ion{H}{i} gas within the framework of our model, one needs to know the $L/\\ell$ ratio. This parameter can be found from the analysis of saturated and optically thin lines of the same ion, since the latter are less affected by the correlation effects. For our case, observations of the \\ion{Ar}{i}\\,$\\lambda1066$ line would be of particular interest since its oscillator strength $f_{1066} = 1/4\\,f_{1048}$. The \\ion{Ar}{i}\\,$\\lambda1066$ line was observed in the $z = 3.4$ DLA system toward the quasar Q0000-2620 and its analysis showed a remarkably similar abundance to the other $\\alpha$-chain elements O, S, and Si, $Z/Z_\\odot \\simeq 1/80$ (Molaro et al. 2000). The similarity of both DLA systems is also supported by the very low molecular hydrogen contents found in I~Zw~18, $f({\\rm H}_2) \\equiv 2N_{{\\rm H}_2}/N_{\\ion{H}{i}} \\ll 10^{-6}$ (Vidal-Madjar et al. 2000), and at $z=3.4$, $f({\\rm H}_2) \\simeq 4\\times10^{-8}$ (Levshakov et al. 2000).  Since Ar can hardly be depleted onto dust grains, but can be partially ionized by nearby UV stellar radiation with energies $h\\nu > 15.76$ eV, the relative abundance of \\ion{Ar}{i} is a good indicator of the intensity of the local photoionizing flux (Sofia \\& Jenkins 1998). Our measurements rule out the presence of a significant amount of partially ionized gas in the damped L$\\alpha$ system I~Zw~18 and, hence, metal abundances in the neutral gas do not require ionization corrections as suggested by Izotov et al. (2000).  Thus, it is extremely important in future observations to investigate low-ion lines to undestand better the kinematic characteristics of the neutral gas bulk motion which will enable us to obtain more reliable estimations of metallicities."
    },
    "0011/astro-ph0011455_arXiv.txt": {
        "abstract": "We present BeppoSAX observations of the two FR~I type radio galaxies Hercules A  (3C 348) and  MRC 0625-536 in the energy range $0.2 - 200$ keV. Data analysis  shows that the X-ray flux from Hercules A is consistent with a diffuse thermal plasma emitting at $T \\approx 4-5$ keV with a possible, but somewhat uncertain, contribution of a softer component at $T \\approx 3 $ keV. The non thermal  emission from the active nucleus must be significantly smaller than the thermal  one, and no indication of relevant core obscuration by a surrounding torus  was detected. The flux from MRC 0625-536 originates form an extended region and has been fitted to a thermal law with $T \\approx 5.7$ keV and with a column density consistent with the galactic absorption. A spatially resolved spectral analysis does not show a relevant variation of the temperature and the metallicity across the diffuse emission zone.  A non thermal spectral component, related to the nuclear activity, may be  present in the innermost region with some possible amount of local obscuration, contributing $\\lapp \\; 10\\%$ to the total luminosity. Hard X-ray emission from  MRC 0625-536 has been detected in the PDS (15 - 200 keV) that may be related  either to its galactic core or to the intracluster region. ",
        "introduction": "The two sources Hercules A (3C 348) and MRC 0625-536 are radio galaxies classified as Fanaroff-Riley I (FR~I) type objects, although Hercules A has indeed the morphology of a FR~I radio galaxy, but with a much higher  radio luminosity, typical of a FR~II object.  The galaxy associated with MRC 0625-536 is the main member of the cluster A 3391, while  Hercules A lies in a group.  Recent spectral and morphological analyses of X-ray data, collected by  ROSAT and ASCA (Otani et al. 1998, Siebert et al. 1999, Gliozzi et al. 2000), have provided  information on the physical and dynamical properties of the environment surrounding these objects and of their active nuclei. Both radio galaxies are embedded in diffuse clouds of X-ray emitting hot plasma of size larger than the extended, radio emitting, regions. The active nucleus may also contribute to the X-ray flux, as observed in several FR~I radio galaxies, and this emission is likely related to the inner region of the relativistic jet (Worrall \\& Birkinshaw 1994, Trussoni et al. 1999, Hardcastle \\& Worrall 1999).  In this paper we present  BeppoSAX observations of Hercules A and MRC 0625-536. This analysis provides new information on the structure of these radio galaxies and their associated clusters. In fact, the instrumentation on board of BeppoSAX allows us to collect data, during the same pointing, that range from the soft band up to $\\approx 200$ keV. We will also examine the variation of the spectral parameters of the X-ray emission across the clusters, relying upon the smoother (with respect to the ASCA detectors) radial behavior of the Point Spread Function of the BeppoSAX MECS detectors.  In Section 2 we summarize the known properties of the two sources, in Section 3 we outline the details of the observations and of the data analysis while in Section 4 we report the results of the spectroscopic analysis. Their implications on the structure of the two sources are discussed in detail in Section 5,  taking into account also previous observations, and in Section 6 we summarize our results. A value of $H_0= 50$ km s$^{-1}$ Mpc$^{-1}$ is assumed throughout. ",
        "conclusions": "We summarize the main results obtained from the BeppoSAX observations and their implications on the properties of these two radio galaxies and their environment.  \\noindent {\\it Hercules  A.} No relevant variation of temperature and metallicity occurs between the inner and outer regions of the cluster. However a plasma with two different temperatures seems to be present, perhaps related to the interaction of the radio components with intracluster gas. It must be borne in mind anyway that the structure of this object  is quite peculiar: at optical wavelengths it is  classified as a group but its X-ray properties are typical of clusters.  The non thermal emission from the active nucleus appears unabsorbed, as expected for this class of radio galaxies, but brighter than expected by its core radio emission. This could be related to the peculiar nature of Hercules A that shows properties of both FR~I and FR~II radio galaxies.  \\noindent {\\it MRC 0625-36.} The X-ray luminosity and the values of $T$ and $\\mu$ are consistent with those expected for this kind of clusters. The lack of strong variation of the temperature and metallicity  across the emitting region suggests that no cooling flow is present, while it could be related to a recent merging process. The nuclear non thermal emission detected in the MECS and in previous  missions is consistent with that expected from radio data,  without excluding however some amount of local obscuration.  Concerning finally the high energy flux detected in the PDS, it seems unlikely that it is related to the cluster or to the radio lobes. On the other hand, a nuclear origin for this emission would imply a luminosity of some units of $10^{44}$ erg s$^{-1}$ in the hard band. Either this emission is completely independent of the flux detected at lower energies, or, to fit the ROSAT, ASCA and BeppoSAX (MECS) data, it must have an very flat spectrum."
    },
    "0011/astro-ph0011349_arXiv.txt": {
        "abstract": "We present [N~{\\sc ii}] and H$\\alpha$ images and high resolution long--slit spectra of the planetary nebula IC\\,4846, which reveal, for the first time, its complex structure and the existence of collimated outflows. The object consists of a moderately elongated shell, two (and probably three) pairs of collimated bipolar outflows at different orientations, and an attached circular shell. One of the collimated pairs is constituted by two curved, extended filaments whose properties indicate a high velocity, bipolar precessing jet. A difference of $\\simeq$ 10 km\\,s$^{-1}$ is found between the systemic velocity of the precessing jets and the centroid velocity of the nebula, as recently report for Hu\\,2-1. We propose that this difference is due to orbital motion of the ejection source in a binary central star. The orbital separation of $\\leq$ 30 AU and period $\\leq$ 100 yr estimated for the binary are similar to those in Hu\\,2-1, linking the central stars of both planetary nebulae to interacting binaries. Extraordinary similarities also exist between IC\\,4846 and the bewildering planetary nebula NGC\\,6543, suggesting a similar formation history for both objects. ",
        "introduction": "IC\\,4846 (PN\\,G\\,027.6$-$09.6) is a compact planetary nebula (PN) whose morphology has not been studied in detail yet. The only available information on its structure is provided by the VLA 6\\,cm continuum observations by Kwok (1985, see also Aaquist \\& Kwok 1990), showing several knots embedded in a faint elongated structure of $\\simeq$ 3$\\times$2 arcsec$^2$ in size. The H$\\beta$ surface brightness ($\\log F({\\rm H}\\beta) = -11.34$, Acker et al. 1992) suggests that IC\\,4846 has a high electron density. This is corroborated by the small [S~{\\sc ii}]$\\lambda$6717,$\\lambda$6731 doublet ratio (Barker 1978; Acker et al. 1992) which reaches the limiting ratio for high electron density ($N_{\\rm e} \\gapprox 10^4$~cm$^{-3}$). Both its small angular size and high density indicate that IC\\,4846 is a young PN. The line intensity ratios (Acker et al. 1992) point to a moderate excitation ([N~{\\sc ii}]$\\lambda$6583/H$\\alpha \\simeq 0.07$, [O~{\\sc iii}]$\\lambda$4959/H$\\beta \\simeq 3.7$, and very faint He~{\\sc ii} $\\lambda$4686). {\\it IUE} low-resolution UV spectra of the central star show a strong P-Cygni profile in the N~{\\sc v} line\\footnote{ Based on INES data from the {\\it IUE} satellite.} with a terminal velocity of $\\simeq$ 1300 km~s$^{-1}$. The occurrence of a strong stellar wind in the central stars of PNe is indicative of their early evolutionary stage (Cerruti-Sola \\& Perinotto 1985).  The shaping processes of PNe occur early in their formation. Therefore, the signatures of these processes in the nebula are more recent and more likely to have withstood the fast nebular evolution in young PNe. This turns the observations of young PNe of chief importance in the understanding of their formation mechanisms. Accordingly, IC\\,4846 was included in our observational program of compact young PNe (see Miranda 1999 and references therein), and narrow--band H$\\alpha$ and [N~{\\sc ii}] images and long--slit high--resolution spectra were obtained to study its structure and internal kinematics. The results reported in this paper evidence the existence of bipolar, precessing collimated outflows in IC\\,4846. ",
        "conclusions": "The observations of IC\\,4846 have disclosed a complex structure composed of three different elements: (1) two apparent elongated shells with different orientations, (2) an attached circular shell and (3) three pairs of point-symmetric components at different orientations with typical properties of collimated outflows. In the following we discuss about the nature, models and possible interpretations of these components.  As already mentioned, there is no kinematic evidence for the elliptical shells in the PV maps of the [N~{\\sc ii}] line, although the extended emission observed in the PV maps may be related to these shells. We also note that the major axis of the outer ellipse is only defined by A1--A2. In fact, without A1--A2 the outer ellipse could not be recognized. Images at higher resolution would be necessary to elucidate the precise geometry of these shells. While A1--A2 and B1--B2 can be interpreted as precessing bipolar outflows, the interpretation of the pair C1--C2 is somewhat difficult. Its kinematic properties and the elongation perpendicular to the radial direction suggest that C1--C2 could represent collimated ejections in which the orientation of the ejection axis has remarkably changed over the ejection time. However, we cannot completely rule out that C1--C2 represents a ring-like structure containing two bright opposite regions. The attached circular shell probably represents a previous slow wind ejected by the AGB progenitor of IC\\,4846. It is similar to other attached shells observed in PNe and its formation may be related to the dynamics of ionization fronts (see Mellema 1995 and references therein).  In the scenario that A1--A2 are precessing bipolar jets, we can obtain estimates for the inclination angle and velocity by assuming that the expansion velocity is constant and that the variation of the radial velocity is due to a systematic change of the angle between the observer and the precessing jet along a precession cone with aperture angle $\\simeq$ 5$^{\\circ}$--10$^{\\circ}$. The observed minimum and maximum radial velocities are assumed to be related to the far and rear side of the precession cone, respectively. Using the model for a biconical flow described by Miranda \\& Solf (1992), we obtain an inclination angle for the precession axis of $\\simeq$ 13$^{\\circ}$--26$^{\\circ}$ with respect to the plane of the sky, and a velocity of $\\simeq$ 100--190 km\\,s$^{-1}$. The (mean) kinematic age is $\\simeq$ 200$\\times$D(kpc) yr, where D is the distance to IC\\,4846, which is extremely uncertain with published values between 2 and 11 kpc (see Acker et al. 1992). The images also show that A1--A2 protrude from the circular shell. This result, the increase of the [N~{\\sc ii}]/H$\\alpha$ ratio at larger distances and the decrease of the radial velocity at the tips of A1--A2 suggest an interaction between the jet fronts and the circular shell, in a similar way to that reported in, e.g., NGC\\,6891 (Guerrero et al. 2000).  It is worthwhile to emphasize the extraordinary similarity between IC\\,4846 and the well-known PN NGC\\,6543 (see Miranda \\& Solf 1992 and Reed et al. 1999). The following structural components are common to these two PNe (we will use for NGC\\,6543 the nomenclature of Miranda \\& Solf 1992): \\begin{enumerate} \\item {\\em Basic elongated shell(s).} In NGC\\,6543 two shells are observed, an outer bipolar shell with a bright ring between the lobes and an inner elliptical one. In IC\\,4846 two elongated shells are suggested by the images (but see above). The possible existence of an equatorial ring should be confirmed by new data. \\item {\\em Bipolar knots at the tips of the shell.} The knots FF' in NGC\\,6543 correspond to B1--B2 in IC\\,4846. \\item {\\em Precessing bipolar jets.} The bipolar jets JJ' in NGC\\,6543 correspond to A1--A2 in IC\\,4846. The similarity between JJ' and A1--A2 is extraordinary. Both appear off-axis with respect to the orientation of the basic shell and are curved, with a change of orientation of 5$^{\\circ}$--10$^{\\circ}$. The radial velocity presents an identical variation along the jets with a steady increase up to a maximum radial velocity and a small decrease at their tips. Finally, an expansion velocity of $\\simeq$ 100--200 km\\,s$^{-1}$ is obtained in both cases. \\item {\\em Point-symmetric components perpendicular to the radial direction.} C1--C2 in IC\\,4846 could correspond to DD' in NGC\\,6543. \\end{enumerate}  There are some differences between both objects, but these are most likely attributable to different orientations, ages and distances. Whatever these differences, the conspicuous kinematic and structural similarities lead us to conclude that NGC\\,6543 and IC\\,4846 have shared a similar formation history.  An interesting result in IC\\,4846 is the systemic velocity shifts among the different components and the centroid of the nebula. In the case of B1--B2 and C1--C2, the shifts are only slightly larger than the relative errors and we cannot conclude whether they are real. In the case of A1--A2, however, the shift is noticeable and we consider it real. Both the remarkable spatial and velocity symmetry observed in A1--A2 exclude that this shift may be ascribed to different ejection velocity, changes (deceleration) in the velocity in one of the components of the pair, different ejection angle or a combination of the three. A more plausible explanation is that the systemic velocity of A1--A2 is indeed intrinsically different from that of the other components and, in particular, from that of the emission centroid. Similar shifts in the systemic velocity of collimated outflows and shell have been recently reported in the PN Hu\\,2-1 (Miranda et al. 2000). Furthermore, both PNe share the presence of two systems of collimated outflows at different orientations, one of them projected along the nebular major axis and one off-axis, suggesting a parallel formation history. Moreover, inspection of the kinematic data of NGC\\,6543 (see Miranda \\& Solf 1992) suggests that a difference of $\\sim$ 7 km\\,s$^{-1}$ could exist between the systemic velocity of JJ' and the systemic velocity of the nebula as defined by the rest of components. In the case of Hu\\,2-1, the shifts of systemic velocity have been interpreted as a direct evidence of orbital motion of the ejection source in a binary central star (Miranda et al. 2000). The velocity shifts detected in IC\\,4846 fit also into this framework. If A1--A2 are ejected from a star in a binary system, the space velocity of the A1--A2 will contain an additional velocity component due to the orbital motion, besides to the component due to the own ejection velocity. Accordingly, the systemic velocity of A1--A2 will be shifted from that of the emission centroid of the nebula (see Miranda et al. 2000 for details). We note that the same would apply to B1--B2 and C1--C2 if their observed velocity shifts were real.  The systemic velocity shift of A1--A2 (10 km\\,s$^{-1}$) is very similar to that found in Hu\\,2-1 (8 km\\,s$^{-1}$) and in NGC\\,6543 (7 km\\,s$^{-1}$). Consequently, the orbital parameters of their binary central stars should be rather similar as well, with an orbital separation $\\leq$ 30 AU and an orbital period $\\leq$ 100 yr, thus linking the central stars of these nebulae to symbiotic stars. These results strongly suggest that the central star of IC\\,4846, similarly to that of Hu\\,2-1 and NGC\\,6543, is a binary star in which mass transfer from a red giant leads to the formation of an accretion disk about a secondary (Mastrodemos \\& Morris 1998, 1999), from which collimated ejections may be generated."
    },
    "0011/astro-ph0011380_arXiv.txt": {
        "abstract": "We present VLT FORS1 spectroscopy that temporally resolves the second caustic crossing of the Bulge K giant source of microlensing event \\event, the first time this has been accomplished for several phases of a caustic transit. The $\\sim$1 \\AA\\ H$\\alpha$ equivalent width of the source star increases slightly as the center of the star egresses the caustic and then plummets by 30\\% during the final limb crossing. These changes are not seen in contemporaneous spectra of control stars in the FORS1 slit, but are qualitatively consistent with expectations from stellar atmosphere models as the caustic differentially magnifies different portions of the stellar face of the target.  Observations such as these in a variety of stellar lines are equivalent to atmospheric tomography and are expected to provide a direct test of stellar models. ",
        "introduction": "\\label{intro}  The steep magnification gradient near the caustics\\footnote{\\footnotetext ~~A caustic is a locus of positions in the source plane for which the magnification is formally infinite.} generated by a binary microlens both magnifies and resolves the background star as the sharply peaked lensing ``beam'' sweeps across the source.  The finite size of the source and its wavelength-dependent surface brightness profile leave their signatures on the shape of the resulting microlensing light curve. The largest effect is the broadening and decrease in amplitude of the caustic peaks in the light curve which is related to the size of the source compared to that of the Einstein ring and can be used to measure the relative proper motion of the lens and source (Alcock \\etal\\ 1997, 2000; Albrow \\etal\\ 1999a, 1999b, 2000, 2001a).  The next highest order term is that due to broadband limb darkening: at most wavelengths, the stellar disk is brighter at center than at limb, which steepens the slope of the light curve when the caustic transits the stellar limb.  These effects were predicted several years ago (Schneider \\& Wagoner 1987; Witt \\& Mao 1994; Witt 1995; Bogdanov \\&  Cherepashchuk 1995; Gould \\& Welch 1996) and have now been used to measure wavelength dependent limb-darkening coefficients for three Bulge giants (Albrow \\etal\\ 1999b, 2000, 2001a) and one A dwarf in the Small Magellanic Cloud (Afonso \\etal\\ 2000).  Such measurements are important to stellar atmosphere physics since very few techniques are available to extract reliable limb-darkening coefficients for typical stars other than the Sun.  Time-resolved spectroscopic monitoring with large aperture telescopes now offers the means to go one step further.  Measurements of the changes in the equivalent width and line shape of individual emission and absorption lines in the source star during a caustic crossing are now feasible (Loeb \\& Sasselov 1995; Valls-Gabaud 1998) and can be expected to yield detailed information about the chemothermal structure of the atmosphere as a function of depth (Heyrovsk\\'y, Sasselov \\& Loeb 2000). Starspots and the polarization and kinetic structure of the stellar envelopes may also be detectable with the aid of microlensing caustics (Simmons, Newson \\& Willis 1995; Simmons, Willis \\& Newsam 1995; Igance \\& Hendry 1999, Bryce \\& Hendry 2000, Heyrovsk\\'y \\& Sasselov 2000).  Fold (line) caustics, such as the one that transited the source star of \\event, would be expected to produce somewhat smaller differential signals ($\\sim$1\\%) than direct transits by (single-lens) point caustics (Gaudi \\& Gould 1999) because they have a broader ``beam pattern''. On the other hand, fold caustic transits are observed much more frequently than are direct point transits.  Detection of the differential signal in individual spectral lines thus requires not only large apertures and efficient spectrographs to achieve the high S/N, but also control or understanding of systematics at the $\\sim$1\\% level.  Following a real-time electronic microlensing alert issued by the EROS group, the PLANET collaboration (Albrow \\etal\\ 1998) commenced intense photometric monitoring of the event \\event.  A secondary alert issued by the MPS collaboration announced an unexpected, sudden, and significant brightening of the object; PLANET increased its sampling rate immediately.  The resulting dense coverage allowed characterization of the first caustic crossing of this binary event by a technique advocated by Albrow \\etal\\ (1999c). Together with subsequent high precision photometry in the caustic trough, a prediction for the timing of the second caustic crossing could then be made, which was announced electronically by PLANET\\footnote{\\footnotetext ~~See http://www.astro.rug.nl/$\\sim$planet/EB2K005.html} a few days in advance, facilitating the VLT observations reported here. The exceptionally long duration of this crossing (due to the glancing angle at which the source left the caustic region) allowed PLANET spectroscopic monitoring from Paranal to be spread over several nights, during which time the source was differentially resolved as the caustic passed over the center and then the limb of the exiting source.  In this Letter, we describe the changes in the equivalent width of the H$\\alpha$ line observed during our VLT FORS1 spectroscopic monitoring of the K-giant source star of \\event, the first time that full temporal coverage has been obtained of a caustic crossing at such high spectral resolution.  We show that these changes in H$\\alpha$ are qualitatively consistent with those expected from stellar atmosphere models for a star of this type; more detailed comparisons with theory over the complete spectrum will be presented elsewhere (Albrow \\etal\\ 2001b). ",
        "conclusions": "Center-to-limb variations are expected in many strong absorption lines of cool giants, including H$\\alpha$.  Such lines are expected to be deeper just before and during the egress of the center of a giant microlensed star from a fold caustic.  Since the line is in emission at the coolest portions of the outer atmosphere, the observed line should become shallower (smaller equivalent width) as the limb egresses and becomes magnified differentially with respect to the rest of the star.  This behavior is confirmed by our observations of \\event\\ throughout the final transit of the K giant source star by the caustic.  Previous attempts to measure spectral line changes during caustic crossings have been encouraging, but hampered by the lack of temporal coverage during those phases of the caustic transit that most significantly resolve source structure. The real-time alert provided by the MACHO team (Alcock \\etal\\ 2000) allowed Lennon \\etal\\ (1996) to take NTT spectra over the peak of the fold crossing in MACHO~96-BLG-3, but these did not extend far enough temporally to detect spectral differences due to source resolution. A difference in H$\\alpha$ equivalent width during and after the grazing point-caustic transit of the M giant source of event MACHO~95-BLG-30 was reported by Alcock et al.\\ (1997), and appears to be qualitatively compatible with 1-D stellar models (Sasselov 1998).  The reported changes in TiO equivalent width for this event (Alcock \\etal\\ 1997) are difficult to interpret as they are as large in the several days after the crossing as during and just after the transit.  Our spectroscopic monitoring of \\event\\ with FORS1 on the VLT is the first to be carried out through all phases of a microlensing caustic crossing. Significant changes in the equivalent width of H$\\alpha$ were detected throughout the transit.  A fiducial K giant model atmosphere convolved with simple fold caustic magnification models exhibits changes in equivalent width that are in qualitative agreement with the observed changes. At the time of this writing, PLANET continues photometric monitoring of \\event\\ which, although far past the caustic structure, is still declining in brightness.  Modeling indicates that the source star passed quite close to a caustic cusp during the second crossing, altering the temporal and spatial behavior of the microlensing magnification ``beam''. A quantitative confrontation with atmospheric models requires proper treatment of the effect of the cusp over the whole stellar disk, and is now in progress (Albrow \\etal\\ 2001b).  Stellar tomography using microlensing caustics is now feasible; adequate alert mechanisms combined with modeling and dense, precise follow-up photometry enable prediction of caustic transits days in advance. Spectra of the source and field control stars of \\event\\ demonstrate that equivalent width changes as small as 0.1\\AA\\ and relative spectral differences of $\\sim$1\\% per resolution element can be observed in two hours (per temporal sampling element) with FORS on the VLT for microlensed Bulge sources. These characteristics are well-matched to monitoring the equivalent width of many lines during a (typical) 8-24 hour caustic transit.  The evolution of spectral line shape, an even more sensitive indicator of stellar atmosphere conditions, will require the higher resolution provided by Echelle instruments. If large aperture telescopes with such spectroscopic ability, especially those in the southern hemisphere, can be flexibly scheduled in real-time, microlensing can meaningfully test theories of stellar atmospheres in the next few years."
    },
    "0011/hep-ph0011160_arXiv.txt": {
        "abstract": "We present a fully nonlinear study of the development of equilibrium after preheating. Preheating is the exponentially rapid transfer of energy from the nearly homogeneous inflaton field to fluctuations of other fields and/or the inflaton itself. This rapid transfer leaves these fields in a highly nonthermal state with energy concentrated in infrared modes. We have performed lattice simulations of the evolution of interacting scalar fields during and after preheating for a variety of inflationary models. We have formulated a set of generic rules that govern the thermalization process in all of these models. Notably, we see that once one of the  fields is amplified through parametric resonance or other mechanisms it rapidly excites other coupled fields to exponentially large occupation numbers. These fields quickly acquire nearly thermal spectra in the infrared, which gradually propagates into higher momenta. Prior to the formation of total equilibrium, the excited fields group into subsets with almost identical characteristics (e.g. group effective temperature). The way fields form into these groups and the properties of the groups depend on the couplings between them. We also studied the onset of chaos after preheating by calculating the Lyapunov exponent of the scalar fields. ",
        "introduction": "The theory of inflation has been highly successful in explaining many of the initial conditions for the Hot Big Bang model as well as providing a mechanism by which the seeds of large scale structure were formed. Typical models of inflation are based on the slow-roll evolution of the homogeneous inflaton scalar field(s) $\\phi$. Inflation ends when the slow-roll regime is dynamically terminated and the field(s) begins to oscillate around the minimum of its effective potential $V(\\phi)$ as in chaotic inflation \\cite{linde}, or ``waterfalls'' towards the minimum of $V$ as in hybrid inflation \\cite{hybr}. After inflation the homogeneous inflaton field(s) decays due to its interactions with other fields or its self-interaction.  If the inflaton decay into other fields were slow as in perturbation theory the created particles would settle into thermal equilibrium as the inflaton decayed. However, the decay of the inflaton typically occurs via rapid, non-perturbative mechanisms collectively known as preheating \\cite{KLSearly}.  The character of preheating may vary from model to model, e.g. parametric excitation in chaotic inflation \\cite{KLS97} and another, specific type of preheating  in hybrid inflation \\cite{hh}, but its distinct feature remains the same: rapid amplification of one or more bosonic fields to exponentially large occupation numbers. This amplification is eventually shut down by backreaction of the produced fluctuations. The end result of the process is a turbulent medium of coupled, inhomogeneous, classical waves far from equilibrium \\cite{KT}.  Despite the development of our understanding of preheating after inflation, the transition from this stage to a  hot Friedmann universe in thermal equilibrium has remained relatively poorly understood. A theory of the thermalization of the fields generated from preheating is necessary to bridge the gap between inflation and the Hot Big Bang.  The details of this thermalization stage depend on the constituents of the fundamental Lagrangian ${\\cal L}(\\phi_i, \\chi_i, \\psi_i, A_{\\mu}, h_{\\mu\\nu}, ...)$ and their couplings, so at first glance it would seem that a description of this process would have to be strongly model-dependent. We have found, however, that many features of this stage seem to hold generically across a wide spectrum of models. This fact is understandable because the conditions at the end of preheating are generally not qualitatively sensitive to the details of inflation. Indeed, at the end of preheating and beginning of the turbulent stage (denoted by $t_*$), the fields are out of equilibrium. We have examined many models and found that at $t_*$ there is not much trace of the linear stage of preheating and conditions at $t_*$ are not qualitatively sensitive to the details of inflation. We therefore expect that this second, highly nonlinear, turbulent stage of preheating may exhibit some universal, model-independent features.  Although a realistic model would include one or more Higgs-Yang-Mills sectors, we treat the simpler case of interacting scalars.  Within this context, however, we consider a number of different models including several chaotic and hybrid inflation scenarios with a variety of couplings between the inflaton and other matter fields.  There are many questions about the thermalization process that we set out to answer in our work. Could the turbulent waves that arise after preheating be described by the theory of (transient) Kolmogorov turbulence or would they directly approach thermal equilibrium? Could the relaxation time towards equilibrium be described by the naive estimate $\\tau \\sim (n \\sigma_{int})^{-1}$, where $n$ is a density of scalar particles and $\\sigma_{int}$ is a cross-section of their interaction? If the inflaton $\\phi$ were decaying into a field $\\chi$, what effect would the presence of a decay channel $\\sigma$ for the $\\chi$ field have on the thermalization process? For that matter, would the presence of $\\sigma$ significantly alter the preheating of $\\chi$ itself, or even destroy it as suggested in \\cite{PR}? How strongly model-dependent is the process of thermalization; are there any universal features across different models? Finally there's the question of chaos. It is known that Higgs-Yang-Mills systems display chaotic dynamics during thermalization \\cite{Biro}. The possibility of chaos in the case of a single, self-interacting inflaton was mentioned in passing in \\cite{KT}, but when we began our work it was unclear at what stage of preheating chaos might appear, and in what way.  Because the systems we are studying involve strong, nonlinear interactions far from thermal equilibrium it is not possible to solve the equations of motion using linear analysis in Fourier space. Instead we solve the scalar field equations of motion directly in position space using lattice simulations. These simulations automatically take into account all nonlinear effects of scattering and backreaction.  Using these numerical results we have been able to formulate a set of empirical rules that seem to govern thermalization after inflation.  These rules qualitatively describe thermalization in a wide variety of models. The features of this process are in some cases very different from our initial expectations.  Section \\ref{setting} gives a brief review of preheating in different inflationary models. This review should serve to motivate our study and place it in the broader context of inflationary cosmology. Sections \\ref{results} and \\ref{others} describe the results of our numerical calculations. Section \\ref{results} describes one simple chaotic inflation model that we chose to focus on as a clear illustration of our results, while section \\ref{others} discusses how the thermalization process occurs in a variety of other models. Section \\ref{chaos} describes the onset of chaos during preheating and includes a discussion of the measurement and interpretation of the Lyapunov exponent in this context. Section \\ref{rules} contains a list of empirical rules that we have formulated to describe thermalization after preheating. Section \\ref{discussion} discusses these results and other aspects of non-equilibrium scalar field dynamics. Finally, there is an appendix that describes our lattice simulations. ",
        "conclusions": "\\label{discussion}  We investigated the  dynamics of interacting scalar fields during post-inflationary preheating and the development of equilibrium immediately after preheating. We used three dimensional lattice simulations to solve the non-linear equations of motion of the classical fields.  There are a number of problems both from the point of view of realistic models of early universe preheating and from the point of view of non-equilibrium quantum field theory that we have not so far addressed. In this section we shall discuss some of them.  Although we considered a series of models of inflation and interactions, we mostly restricted ourselves to four-legs interactions. (The sole exception was the hybrid inflation model, which develops a three-legs interaction after symmetry breaking.) This meant we still had a residual homogeneous or inhomogeneous inflaton field. In realistic models of inflation and preheating we expect the complete decay of the inflaton field. (There are radical suggestions to use the residuals of the inflaton oscillations as dark matter or quintessence, but these require a great deal of fine tuning.) The problem of residual inflaton oscillations can be easily cured by three-legs interactions. In the scalar sector three-legs interactions of the type $g^2 v \\phi \\chi^2$ may result in stronger preheating. Yukawa couplings $h \\bar \\psi \\phi \\psi$ will lead to parametric excitations of fermions \\cite{GK98}.  There are subtle theoretical issues related to the development of precise thermal equilibrium in quantum and classical field theory due to the large number of degrees of freedom, see e.g. \\cite{approx}. In our simulations we see the flattening of the particle spectra $n_k$ and we describe this as an approach to thermal equilibrium, but in light of these subtleties we should clarify that we mean {\\it approximate} thermal equilibrium.  Often classical scalar fields in the kinetic regime display transient Kolmogorov turbulence, with a cascade towards both infrared and ultraviolet modes \\cite{Cond,waves}. In our systems it appears that the flux towards ultraviolet modes is occurring in such a way as to bring the fields closer to thermal equilibrium (\\ref{wave}). Indeed, the slope of the spectra $n_k$ at the end of our simulations is close to $-1$. However, given the size of the box in these simulations we can little say about the phase space flux in the direction of infrared modes. This question could be addressed, for example, with the complementary method of chains of interacting oscillators, see \\cite{Cond}. This is an interesting problem because an out-of-equilibrium bose-system of interacting scalars with a conserved number of particles can, in principle, develop a bose-condensate. It would be interesting to see how the formation of this condensate would or would not take place in the context of preheating in an expanding universe. One highly speculative possibility is that a cosmological bose condensate could play the role of a late-time cosmological constant.  The highlights of our study for early universe phenomenology are the following. The mechanism of preheating after inflation is rather robust and works for  many different systems of interacting scalars. There is a stage of turbulent classical waves where the initial conditions for preheating are erased. Initially, before all the fields have settled into equilibrium with a uniform temperature, the reheating temperature may be different in different subgroups of fields. The nature of these groupings is determined by the coupling strengths."
    },
    "0011/astro-ph0011394_arXiv.txt": {
        "abstract": "We have quantified the galaxy environments around a sample of $0.5 \\leq z \\leq 0.8$ radio-quiet quasars using the amplitude of the spatial galaxy--quasar correlation function, $B_{\\rm gq}$. The quasars exist in a wide variety of environments, some sources are located in clusters as rich as Abell class 1--2 clusters, whereas others exist in environments comparable to the field. We find that on average, the quasars prefer poorer clusters of $\\approx$ Abell class 0, which suggests that quasars are biased tracers of mass compared to galaxies. The mean $B_{\\rm gq}$ for the sample is found to be indistinguishable from the mean amplitude for a sample of radio-loud quasars matched in redshift and optical luminosity. These observations are consistent with recent studies of the hosts of radio-quiet quasars at low to intermediate redshifts, and suggest that the mechanism for the production of powerful radio jets in radio-loud quasars is controlled by processes deep within the active galactic nucleus itself, and is unrelated to the nature of the hosts or their environments. ",
        "introduction": "\\label{section:s1}  \\footnotetext{\\footnotemark[4]Email: wold@astro.su.se}  The study and interpretation of the differences and similarities between radio-loud and radio-quiet quasars (RLQs and RQQs, hereafter) has kept astronomers busy for a long time. The spectral energy distributions of RLQs and RQQs differ only in points of detail at all wavelength bands, apart from the radio. Both quasar types possess broad emission lines and have a stellar appearance in even the highest resolution optical images. They also show the same evolution in number density with a peak at $z\\sim 2$--3, but the RQQs outnumber the RLQs by a factor of 10. RQQs have radio luminosities 2--3 orders of magnitude lower than their radio-loud counterparts, and on radio maps they are compact with typically only a weak, flat-spectrum component coincident with the optical nucleus. On the other hand, the RLQs have extended lobes of radio emission with hotspots at the outer edges of the radio structure. The radio-emitting lobes are fed by powerful jets emerging from a bright, central core, but RQQs can also have jet-like structures (Blundell et al.\\ 1996; Blundell \\& Beasley 1998), although with bulk kinetic powers $\\sim 10^{3}$ times lower than for RLQs (Miller, Rawlings \\& Saunders 1993). This suggests that both quasar populations have jet-producing central engines, but that the efficiency of the jet production mechanism is very different in the two cases.  One way to learn more about how the two quasar types relate, is to study their host galaxies and their galactic environment. A longstanding belief that RQQs are hosted by spiral galaxies and RLQs by ellipticals is now being questioned. Recent studies (e.g.\\ Dunlop et al.\\ 1993; McLeod \\& Rieke 1994; Taylor et al.\\ 1996; Bahcall et al.\\ 1997; Boyce et al.\\ 1998; McLure et al.\\ 1999; Hughes et al.\\ 2000) have found that powerful quasars at $z \\la 0.5$, both radio-loud and radio-quiet, seem to exist in galaxies with luminosities $>L^{*}$, the luminosity at the break in the luminosity function. Still, a clear picture has not emerged, with some studies claiming a high fraction of disk morphologies amongst the radio-quiets (e.g.\\ Percival et al.\\ 2000), whilst others suggest that nearly all quasars are in giant ellipticals (e.g.\\ McLure et al.\\ 2000). Besides the technical difficulties associated with studying host galaxies, quasar optical luminosity and selection effects may affect the results. E.g.\\ the perhaps somewhat artificial distinction between Seyfert galaxies and quasars at $M_{V}=-23$ leads to confusion as to whether an active galactic nucleus (AGN) is a quasar or a Seyfert. Seyfert galaxies are spirals, are found at lower redshifts and have lower AGN luminosities than RLQs, so very few RLQs exist for comparison. Selection effects may therefore have caused spirals to be preferred as hosts of RQQs and giant ellipticals to be preferred as hosts of RLQs.  Quasar environments have also been investigated on much larger scales than those of typical host galaxies. RLQs at $0.5 \\la z \\la 0.8$ are often found to be associated with groups or clusters of galaxies (Yee \\& Green 1987; Ellingson, Yee \\& Green 1991a; Wold et al.\\ 2000 (hereafter Paper~I)), with richnesses varying from poor groups to Abell Class 1 clusters or richer. This is perhaps not surprising since giant elliptical galaxies, found to be the hosts of powerful RLQs, frequently reside in the centres of galaxy clusters. But what is the galaxy environment like for RQQs in this redshift range, and how does it compare with the environment around RLQs with comparable AGN luminosities?  In this paper, we try to answer this question by investigating a sample of RQQs matched in redshift and optical luminosity to a sample of RLQs analysed in Paper~I. We quantify the galaxy environment using the amplitude of the spatial galaxy--quasar cross-correlation function, $B_{\\rm gq}$, often referred to as the `clustering amplitude'. The first attempt at comparing the environments around RLQs and RQQs using the clustering amplitude was made by Yee \\& Green \\shortcite{yg84}, who found only a marginal difference. Later, Yee \\& Green \\shortcite{yg87} did deeper imaging of the same RQQ fields, and added more RLQ fields to the sample. They repeated the analysis with improved parameters, and obtained an even smaller difference, but due to the small number of RQQs (seven) in their sample, they were unable to draw any firm conclusions.  The problem was later addressed by Ellingson et al.\\  \\shortcite{eyg91a}, who also used the clustering amplitude to quantify the Mpc-scale environment around a sample of 65 $17.5 < m < 19$ quasars at $0.3 < z < 0.6$, of which half were radio-quiet. They found that the RLQs exist more often in rich galaxy environments than the RQQs do, which was believed to support the RQQ/spiral host--RLQ/elliptical host hypothesis. Since this study, there has been little work aimed at comparing the environments of the two quasar populations at moderate redshifts, but at both lower and higher redshifts several studies indicate that contrary to the conclusion reached by Ellingson et al., the galaxy environment around RLQs and RQQs is similar. But also here, a clear picture has not yet emerged, since several other studies seem to support Ellingson et al's conclusion.  Some investigators use the clustering amplitude, which makes comparison between studies similar to ours easier, whereas other studies have other aims, and therefore use different techniques and methods. These studies can basically be divided into two categories. The first category includes measurements of the {\\em angular} galaxy--quasar cross-correlation function by cross-correlating quasar catalogues with galaxy counts. The aim is to establish whether optically and X-ray selected quasars (i.e.\\ mostly RQQs) can be used as unbiased tracers of galaxies. This is important to know, since the large quasars surveys currently in progress, like the 2dF \\cite{shanks00} and the Sloan Digital Sky Survey \\cite{fan99}, will be used to study large-scale structure out to high redshifts, but require knowledge about how quasars are biased with respect to galaxies. The second category includes more or less qualitative searches for galaxies in quasar and radio galaxy fields at both high and low redshifts. E.g.\\ data obtained for host galaxy studies, at typically $z \\la 0.3$, can be utilized for environment studies, but the high resolution imaging required to study hosts often results in a small field of view and therefore a relatively small sampling radius.  Investigations of the galaxy--quasar angular correlation function are often concerned with scales larger than typical cluster scales. At low redshifts, Smith, Boyle \\& Maddox \\shortcite{sbm95} selected 169 ($z<0.3$) optically identified quasars from the EMSS (Einstein Medium Sensitivity Survey), and cross-correlated their positions with $B_{J}<20.5$ galaxies from the APM galaxy catalogues in a region of radius 0\\fdg5 centered on the quasar. They found that the amplitude of the angular galaxy--quasar cross-correlation function was identical to that of the APM galaxy angular correlation function, implying that the quasars inhabit regions of space similar to that of normal galaxies. (At $z=0.3$, 0\\fdg5 corresponds to almost 10 Mpc). In a follow-up study, Smith, Boyle \\& Maddox \\shortcite{sbm00} cross-correlated X-ray selected quasars at $0.3 \\leq z \\leq 0.7$ with $V\\leq23$ galaxies on CCD images. They found a 2$\\sigma$ excess around the $z \\leq 0.5$ quasars on scales $<1$ arcmin, but no excess for the $z>0.5$ quasars. As an overall result, they claimed that the galaxy--quasar correlation function is consistent with that for faint galaxies. The quasars in their sample span a wide range in luminosity, from $M_{V} \\approx -27$ to $-$20 ($H_{0}=50$ km\\,s$^{-1}$\\,Mpc$^{-1}$), but no difference between high and low luminosity sub-samples was observed.  Croom \\& Shanks \\shortcite{cs99} cross-correlated 150 quasars selected both in the optical and in X-rays with $b_{J}<23$ galaxies on AAT (Anglo-Australian Telescope) photographic plates. They found the angular cross-correlation to be marginally negative, indicative of a small anti-correlation, possibly due to gravitational lensing. Most of the quasars in their sample lie at redshifts 1.0--2.2, and they claimed to reach $M^{*}$ at $z\\approx1$ and $M^{*}-1$ at $z\\approx1.5$. Possibly, this is not deep enough, as even $M^{*}$ might be too shallow into the luminosity function to detect excess galaxies at $z=1.0$--2.2 \\cite{ylc99}.  Searches for galaxies in quasar and radio galaxy fields in connection with host studies have been conducted by e.g.\\ Smith \\& Heckman \\shortcite{sh90} who generated catalogues of companion galaxies within 100$h^{-1}$ Mpc of 31 $z < 0.3$ quasars and 35 powerful radio galaxies. They found that RQQs only have half the number of nearby, bright companions as RLQs and radio galaxies at the 95 per cent confidence level. Somewhat inconsistent with Smith \\& Heckman's result is the result obtained by Fisher et al.\\ \\shortcite{fisher96} who found no significant difference in the environments around optically luminous RLQs (six) and RQQs (14) at $z \\la 0.3$. Supporting Fisher et al's result is the recent study by McLure \\& Dunlop \\shortcite{md00}, who quantified the environments around 44 powerful, $z\\approx0.2$ AGN, of which 13 were RLQs and 21 were RQQs (their sample includes Fisher et al's quasars). They found the mean clustering amplitude for the two quasar samples to be consistent with each other, and corresponding to Abell richness class 0 clusters. Dunlop et al.\\ \\shortcite{dunlop93}, studying quasar hosts in the near-infrared, also discuss companion galaxies. Their sample consisted of 14 RLQs and 14 RQQs at $z<0.35$ matched in redshift and $V$ magnitude. They found that the surface density of bright ($K > 18$) companion galaxies was a factor of $\\approx 2.5$ higher than in the field, and the excess was similar for both RLQs and RQQs.  The high redshift studies of quasar environments typically concentrate at $z \\ga 1$. Quasars (and radio galaxies) are used as pointers to high redshift galaxies in order to learn about galaxy formation and evolution and to search for high redshift (proto)clusters. Galaxies, and even clusters, are certainly found at $z \\ga 1$ using this method (e.g.\\ Dickinson 1994; Deltorn et al.\\ 1997; Hall \\& Green 1998; Tanaka et al.\\ 2000, Chapman, McCarthy \\& Persson 2000; Best 2000).  Hintzen, Romanishin \\& Valdes \\shortcite{hrv91} and Boyle \\& Couch \\shortcite{bc93} investigated $0.9 < z < 1.5$ RLQs and RQQs, respectively. Hintzen et al.\\ found a significant galaxy excess within a 15 arcsec radius ($\\approx160$ kpc) centered on the quasars in their sample, whereas Boyle \\& Couch \\shortcite{bc93} found no evidence for excess. However, Boyle \\& Couch's RQQs are 20 times fainter in the optical than the RLQs studied by Hintzen et al., and Boyle \\& Couch discuss this as a possible cause for the discrepancy.  In contrast to this, Hutchings, Crampton \\& Johnson \\shortcite{hcj95} found on average a $>5$$\\sigma$ excess of galaxies around both RLQs and RQQs (14 in total) at $z \\approx 1.1$. They saw no significant difference between RLQ and RQQ environments, which was found to consist of compact groups of starbursting galaxies. Also, Hutchings \\shortcite{hutchings95a}, who looked at galaxy companions to seven $z=2.3$ quasars (only one RLQ) in the $R$-band found an overall $>5$$\\sigma$ excess in a $\\sim 1$ arcmin region around the quasars. However, Hutchings et al.\\ \\shortcite{hutchings99}, using near-infrared imaging of both RLQs and RQQs at $0.9 < z < 4.2$, note that the RQQs occupy poorer environments than the RLQs, but do not draw any firm conclusions due to the small and still incomplete sample. Teplitz, McLean \\& Malkan \\shortcite{tmm99} also report that the number counts of galaxies in RQQ fields at $0.95<z<1.5$ in $J$- and $K'$-band is comparable to the field counts, but their data is not very deep (only reaching $\\approx M^{*}$ at $z \\approx 1$).  In summary, although the results seem to differ, the consensus about hosts and environments of RQQs -- that they are typically late-type galaxies in poor environments -- is being increasingly questioned. Several recent observations at high and low redshifts seem to indicate that RQQs live in richer than average galaxy environment, and powerful RQQs at $z \\la 0.5$ are found in luminous, possibly early-type, galaxies.  Apart from Ellingson et al's \\shortcite{eyg91a} study at $0.3 < z < 0.6$, there have been no attempts to investigate what environments RQQs prefer at intermediate redshifts. We have therefore obtained data on a sample of 20 RQQs at $0.5 \\leq z \\leq 0.8$, and quantified the galaxy richness in the quasar fields using a similar analysis as Ellingson et al. The sample is described in Section~\\ref{section:s2}, where we also discuss the radio loudness of some of the quasars. In Section~\\ref{section:s3}, we describe the observations and our observing strategy. The control fields used in the analysis are also commented upon here, especially how the galaxy counts vary with galactic latitude. The analysis using the clustering amplitude is briefly overviewed in Section~\\ref{section:s4}, before the results are presented in Section~\\ref{section:s5}. In Section~\\ref{section:s51} we examine the colours and the radial distribution of the galaxies in the quasar fields. Using the clustering amplitude allows us to compare our results with some of the above mentioned studies, and this is done in Section~\\ref{section:s6}, and in Section~\\ref{section:s7} we are able to directly and consistently compare with the sample of RLQs that was presented in Paper~I. Finally, we discuss the results in Section~\\ref{section:s8}, and draw the conclusions in Section~\\ref{section:s9}.  Our assumed cosmology has $H_{0}=50$ km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{0}=1$ and $\\Lambda=0$.  \\begin{table*} \\begin{minipage}{13.5cm}  \\caption{The quasar sample and log of observations. All images were obtained using the HiRAC equipped with a 1k SiTe CCD with pixel scale 0.176 arcsec, unless otherwise is specified in footnotes. In the 6th column, we list the FWHM of the seeing in the combined image (in the longest wavelength filter when there are observations in two filters).}  \\begin{tabular}{llllllll} Quasar   & $z$  & Filter & Exp. time  & Date  &  Seeing & Remarks\\\\ &           &   & (s)  &           & (arcsec) & \\\\ &           &     &           &       &    & \\\\ LBQS 0007$-$0003  & 0.698 & $I$ & 4$\\times$600 & 96/07/23 & 1.1 & photometric\\\\  LBQS 0020$-$0300  & 0.580 & $V$, $R$ & 4$\\times$600 & 96/07/25 & 0.8 & dust \\\\  BFSP 12344$-$0021\\footnotemark[1]\\footnotetext{\\footnotemark[1]Imaged using ALFOSC with a 2k Loral CCD with pixel scale 0.189 arcsec.} & 0.753 & $I$ & 5$\\times$300 & 98/04/28 & 0.7 & variable transparency \\\\ &      &      $I$ &   6$\\times$300\\footnotemark[3]  &  00/01/03 &   &  photometric\\\\  BFSP 12355$+$0107\\footnotemark[2]\\footnotetext{\\footnotemark[2]Imaged with the HiRAC camera, but with a 2k Loral CCD with pixel scale 0.11 arcsec.} & 0.720 & $I$ & 9$\\times$300 & 97/05/14 & 0.5 & photometric \\\\  BFSP 12364$-$0053\\footnotemark[1] & 0.727 & $I$ & 8$\\times$300 & 99/05/19,23 & 0.9 & variable transparency\\\\ &         &      $I$ & 1$\\times$300\\footnotemark[3]\\footnotetext{\\footnotemark[3]Exposures taken during photometric conditions in order to calibrate the other images obtained under less photometric conditions.}  &  00/01/03 &   & photometric \\\\  BQS 1333$+$176\\footnotemark[1]  & 0.554 & $R$ & 4$\\times$600 & 00/01/03 & 0.7 & photometric \\\\  BFSP 15152$+$0244\\footnotemark[2] & 0.608 & $R$ & 4$\\times$600 & 97/05/14 & 0.6 & photometric\\\\  BFSP 15196$+$0220   & 0.742 & $I$ & 4$\\times$600 & 96/07/24 & 0.8 & dust \\\\  BFSP 15199$+$0247  & 0.768 & $R$, $I$ & 5$\\times$600 & 96/07/23 & 0.6 & photometric\\\\  BFSP 15202$+$0236  & 0.765 & $R$, $I$ & 4$\\times$600 & 96/07/22 & 0.6 & photometric\\\\  BQS  1538$+$477    & 0.770 & $R$, $I$ & 4$\\times$600 & 96/07/21 & 0.6 & photometric\\\\  BFSP 22011$-$1857 & 0.615 & $V$, $R$ & 4$\\times$600 & 96/07/21 & 1.0 & photometric\\\\  LBQS 2235$+$0054  & 0.529 & $V$, $R$ & 4$\\times$600 & 96/07/21 & 1.0 & photometric\\\\  LBQS 2238$+$0133  & 0.714 & $R$, $I$ & 4$\\times$600 & 96/07/22 & 0.7 & photometric\\\\  LBQS 2239$-$0055 & 0.680 & $R$, $I$ & 4$\\times$600 & 96/07/23 & 0.9 & photometric \\\\  LBQS 2245$-$0055  & 0.801 & $R$, $I$ & 5$\\times$600 & 96/07/24 & 1.3 & dust \\\\  LBQS 2348$+$0210  & 0.504 & $V$, $R$ & 4$\\times$600 & 96/07/24 & 1.3 & dust\\\\  LBQS 2348$+$0148 & 0.749 & $I$ & 4$\\times$600 & 96/07/25 & 0.8 & dust\\\\  LBQS 2350$-$0012  & 0.561 & $V$, $R$ & 4$\\times$600 & 96/07/23 & 0.9 & photometric\\\\  LBQS 2353$-$0153  & 0.672 & $R$, $I$ & 4$\\times$600 & 96/07/22 & 0.6 & photometric\\\\ \\end{tabular} \\end{minipage} \\label{table:table1} \\end{table*}  \\begin{table*} \\begin{minipage}{13.0cm}  \\caption{Radio fluxes at 1.4 GHz from NVSS and FIRST image cutouts. Quasars with radio detections in NVSS or FIRST (or both) are marked with an $R$ in the first column. For the other quasars, we give upper limits that correspond to 3$\\sigma$ detections. Rest frame fluxes at 5 GHz and 2500 {\\AA} are denoted $S_{\\rm 5 GHz}$ and $S_{2500}$, and their ratio is $R_{\\rm 5 GHz}$. The radio power is denoted $P_{\\rm 5 GHz}$, and $\\alpha_r$ is the radio spectral index between 1.4 and 5 GHz.}  \\begin{tabular}{llllllll} Quasar & NVSS  & FIRST & $S_{\\rm 5 GHz}$ &$\\alpha_r$& $S_{\\rm 2500}$ & $R_{\\rm 5 GHz}$ & $P_{\\rm 5 GHz}$ \\\\ & (mJy) & (mJy) & (mJy)           &          & (mJy)          &                 & W\\,Hz$^{-1}$    \\\\ &       &       &                 &          &                &                 &           \\\\ LBQS 0007$-$0003  & $<$ 1.5 & $<$ 0.5 & $<$0.61     & 0.50 & 0.10 & $<$ 6.1 & $<$ 1.6$\\times$10$^{24}$ \\\\  LBQS 0020$-$0300  & $<$ 1.8 & no map  & $<$0.76     & 0.50 & 0.21 & $<$ 3.6   & $<$ 1.4$\\times$10$^{24}$ \\\\  BFSP 12344$-$0021$^{\\rm R}$ & 4.8     & 5.2     & 1.92        & 0.50 & 0.03 & 64.0 & 6.1$\\times$10$^{24}$ \\\\  BFSP 12355$+$0107 & confused& $<$ 0.5 & $<$0.20     & 0.50 & 0.02 & $<$ 10.0  & $<$ 5.8$\\times$10$^{23}$ \\\\  BFSP 12364$-$0053 & $<$ 1.5 & $<$ 0.5 & $<$0.60     & 0.50 & 0.01 & $<$ 60.0  & $<$ 1.8$\\times$10$^{24}$\\\\  BQS 1333$+$176$^{\\rm R}$    & 37.6    & no map  & 18.54       & 0.32\\footnotemark[1]\\footnotetext{\\footnotemark[1]Spectral index found between flux at 1.4 GHz and flux measured at 5 GHz by Kellermann et al.\\ (1989)}  & 0.60 & 30.9 & 3.0$\\times$10$^{25}$ \\\\  BFSP 15152$+$0244 & $<$ 1.5 & $<$ 0.5 & $<$0.63     & 0.50 & 0.02 & $<$ 31.5  & $<$ 1.3$\\times$10$^{24}$ \\\\  BFSP 15196$+$0220 & $<$ 1.5 & $<$ 0.5 & $<$0.60     & 0.50 & 0.03 & $<$ 20.0  & $<$ 1.9$\\times$10$^{24}$ \\\\  BFSP 15199$+$0247 & $<$ 1.5 & $<$ 0.5 & $<$0.60     & 0.50 & 0.03 & $<$ 20.0  & $<$ 2.0$\\times$10$^{24}$ \\\\  BFSP 15202$+$0236$^{\\rm R}$ & $<$ 1.5 & 1.0     & 0.40        & 0.50 & 0.01 & 40.0 & 1.3$\\times$10$^{24}$ \\\\  BQS 1538$+$447$^{\\rm R}$ & 95.2    & 85.4    & 26.28       & 1.02\\footnotemark[1] & 1.00 & 26.3 & 8.8$\\times$10$^{25}$ \\\\  BFSP 22011$-$1857 & $<$ 1.5 & no map  & $<$0.62     & 0.50 & 0.05 & $<$ 12.4  & $<$ 1.3$\\times$10$^{24}$ \\\\  LBQS 2235$+$0054$^{\\rm R}$ & $<$ 1.5 & 0.83    & 0.94        & $-$0.65\\footnotemark[2]\\footnotetext{\\footnotemark[2]Spectral index between 1.4 GHz flux and flux measured at 8.4 GHz by Hooper et al.\\ (1996)} & 0.02 & 47.0 & 1.4$\\times$10$^{24}$ \\\\  LBQS 2238$+$0133  & $<$ 1.5 & no map  & $<$0.61     & 0.50 & 0.22 & $<$ 2.8  & $<$ 1.7$\\times$10$^{24}$ \\\\  LBQS 2239$-$0055  & $<$ 1.5 & $<$ 0.5 & $<$0.61     & 0.50 & 0.12 & $<$ 5.1 & $<$ 1.6$\\times$10$^{24}$ \\\\  LBQS 2245$-$0055$^{\\rm R}$  & $<$ 1.5 & 1.3     & 0.51        & 0.50 & 0.30 & 1.7  & 1.9$\\times$10$^{24}$  \\\\  LBQS 2348$+$0210$^{\\rm R}$  & 41.8    & no map  & 25.05       & 0.12\\footnotemark[2] & 0.17 & 143.4 & 3.3$\\times$10$^{25}$ \\\\  LBQS 2348$+$0148  & $<$ 1.5 & no map  & $<$0.60     & 0.50 & 0.08 & $<$ 7.5  & $<$ 1.9$\\times$10$^{24}$ \\\\  LBQS 2350$-$0012  & $<$ 1.5 & $<$ 0.5 & $<$0.64     & 0.50 & 0.32 & $<$ 2.0  & $<$ 1.1$\\times$10$^{24}$ \\\\  LBQS 2353$-$0153  & $<$ 1.5 & $<$ 0.5 & $<$0.61     & 0.50 & 0.57 & $<$ 1.1 & $<$ 1.5$\\times$10$^{24}$ \\\\ \\end{tabular} \\label{table:table2} \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{section:s9}  The RQQs studied in this paper live in a diversity of environments, from field-like environments to galaxy clusters as rich as Abell class 1--2. On average, the sample has a clustering amplitude corresponding to groups or poorer clusters of Abell class $\\approx 0$, and is {\\em indistinguisable} from the mean amplitude found for a sample of RLQs matched in redshift and AGN luminosity.  Our finding that the average environment around RLQs and RQQs at $0.5 \\leq z \\leq 0.8$ is statistically indistinguishable agrees with what is found for powerful quasars at $z\\approx0.2$ by Fisher et al.\\ \\shortcite{fisher96} and McLure \\& Dunlop \\shortcite{md00}. There is agreement also in terms of the amount of clustering, suggesting that quasar environments, for both RLQs and RQQs, do not undergo much evolution between $z\\approx0.2$ and $\\approx0.8$. These results are consistent with host galaxy studies finding that both powerful RLQs and RQQs live in massive, elliptical galaxies (e.g.\\ Dunlop et al.\\ 1993; Bahcall et al.\\ 1997; McLure et al.\\ 1999). We therefore believe that the consensus that RQQs occupy systematically poorer environments than RLQs should be put under doubt.  Since our RQQs on average appear in cluster environments they are most likely biased with respect to galaxies. Galaxy clusters are also biased (e.g.\\ Lilje \\& Efstathiou 1988; Croft, Dalton \\& Efstathiou 1999), so the RQQs in clusters must have a bias factor that is greater than that for normal galaxies. This suggests that RQQs can not be used as unbiased tracers of galaxies when studying large scale stucture using quasars surveys.  That RLQs and RQQs on average live in equally rich environments implies that the origin of the radio loudness is not linked to the environments on Mpc scales, but may instead be buried within the central parts of the host galaxy. This is consistent with the suggestion that the fuelling of radio jets is linked to accretion onto a central black hole (Rawlins \\& Saunders 1991; Serjeant et al.\\ 1998).  A colour-magnitude diagram of the richest RQQ fields in this study shows hints of a red sequence at the expected colour of $z = 0.7$--0.8 early-type galaxies. With spectroscopy, these fields will likely turn out to contain galaxy clusters at $z=0.7$--0.8. RLQs and radio galaxies are often used to find and study galaxy clusters at higher redshifts, and our result that RLQs and RQQs inhabit similar environment at $z \\sim 0.7$--0.8 should therefore increase the pool of likely cluster candidates at high redshifts. This may be important for theories of galaxy formation, since it allows us to find and study high redshift galaxy clusters that are not selected on the basis of extreme optical or X-ray luminosity."
    },
    "0011/astro-ph0011488_arXiv.txt": {
        "abstract": "We investigate a stationary pair production cascade in the outer magnetosphere of a spinning neutron star. The charge depletion due to a global current, causes a large electric field along the magnetic field lines. Migratory electrons and/or positrons are accelerated by this field to radiate curvature gamma-rays, some of which 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 distribution functions of particles and gamma-rays. If no current is injected at neither of the boundaries of the accelerator, the gap is located around the so-called null surface, where the local Goldreich-Julian charge density vanishes. However, we first find that the gap position shifts outwards (or inwards) when particles are injected at the inner (or outer) boundary. Applying the theory to the Crab pulsar, we demonstrate that the pulsed TeV flux does not exceed the observational upper limit for moderate infrared photon density and that the gap should be located near to or outside of the null surface so that the observed spectrum of pulsed GeV fluxes may be emitted via curvature process. ",
        "introduction": "The EGRET experiment on the Compton Gamma Ray Observatory has detected pulsed signals from seven rotation-powered pulsars (e.g., for Crab, Nolan et al. 1993, Fierro et al. 1998). The modulation of the $\\gamma$-ray light curves at GeV energies testifies to the production of $\\gamma$-ray radiation in the pulsar magnetospheres either at the polar cap (Harding, Tademaru, \\& Esposito 1978; Daugherty \\& Harding 1982, 1996; Sturner, Dermer, \\& Michel 1995; Shibata, Miyazaki, \\& Takahara 1998), or at the vacuum gaps in the outer magnetosphere (Cheng, Ho, \\& Ruderman 1986a,b, hereafter CHR; Chiang \\& Romani 1992, 1994; Romani and Yadigaroglu 1995; Romani 1996; Zhang \\& Cheng 1997, heafter ZC97). Effective $\\gamma$-ray production in a pulsar magnetosphere may be extended to the very high energy (VHE) region above 100 GeV as well; however, the predictions of fluxes by the current models of $\\gamma$-ray pulsars are not sufficiently conclusive (e.g., Cheng 1994). Whether or not the spectra of $\\gamma$-ray pulsars continue up to the VHE region is a question which remains one of the interesting issues of high-energy astrophysics.  In the VHE region, positive detections of radiation at a high confidence level have been reported from the direction of the Crab pulsar (Nel et al. 1993). However, as for {\\it pulsed} TeV radiation, only the upper limits have been, as a rule, obtained. If the VHE emission originates the pulsar magnetosphere, a significant fraction of them can be expected to show pulsation. Therefore, the lack of {\\it pulsed} TeV emissions provides a severe constraint on the modeling of particle acceleration zones in a pulsar magnetosphere.  In fact, in CHR picture, the magnetosphere should be optically thick for pair--production in order to reduce the TeV flux to an unobserved level by absorption. This in turn requires very high luminosities of infrared photons. However, the required IR fluxes are generally orders of magnitude larger than the observed values (Usov 1994). We are therefore motivated by the need to contrive an outer--gap model which produces less TeV emission with a moderate infrared luminosity.  High-energy emission from a pulsar magnetosphere, in fact, crucially depends on the acceleration electric field, $\\Ell$, along the magnetic field lines. It was Hirotani and Shibata (1999a,b,c; hereafter Papers I, II, III), and Hirotani (2000a,b,c; hereafter Papers IV, V, VI) who first solved the spatial distribution of $\\Ell$ together with particle and $\\gamma$-ray distribution functions. They demonstrated that a stationary gap is formed around the null surface at which the local Goldreich--Julian charge density, \\begin{equation} \\rho_{\\rm GJ}= \\frac{\\Omega B_z}{2\\pi c}, \\label{eq:def_rhoGJ} \\end{equation} vanishes, where $B_z$ is the component of the magnetic field along the rotation axis, $\\Omega$ the angular frequency of the neutron star, and $c$ the speed of light. Equation (\\ref{eq:def_rhoGJ}) is valid unless the gap is located close to the light cylinder, of which distance from the rotation axis is given by $ \\rlc= c / \\Omega$. In this letter, we develop the method presented in Paper VI, by investigating the case when particles flow into the gap from the inner or the outer boundaries.  In the next two sections, we describe the physical processes of pair production cascade and the resultant $\\gamma$-ray emission. We then apply the theory to the Crab pulsar and present the expected $\\gamma$-ray spectra in \\S~4. In the final section, we compare the results with ZC97. ",
        "conclusions": "\\label{sec:discussion}  In summary, we have developed a one-dimensional model for an outer-gap accelerator in the magnetosphere of a rotation-powered pulsar. When the electronic current flows into the gap from the outer boundary, the gap shifts inwards to emit very soft GeV emissions. Applying this method to the Crab pulsar, we find that the gap should be located near to or outside of the null surface, so that the observed spectrum of pulsed GeV fluxes may be emitted via curvature process. By virtue of the absorption by the dense IR field in the magnetosphere, the problem of excessive TeV emission does not arise.  Let us briefly compare the present method with ZC97, who considered that the gap width is limited by the surface X-rays due to the bombardment of the particles produced in the gap. The magnetospheric X-rays considered in this paper is much denser than the surface X-rays due to the bombardment. As a result, the localized gap in the present paper produces less intrinsic TeV flux compared with what would be obtained in ZC97 picture.  For cases~1, 2, and 3, the intrinsic TeV luminosity is comparable or less than the GeV one. Therefore, the Lorentz factors are limited primarily by curvature process (eq.[\\ref{eq:terminal}]). For case~4, however, the intrinsic TeV luminosity well exceeds the GeV one; therefore, the radiation-reaction forces are due to IC scatterings rather than the curvature process. In fact, we may expect a sufficient GeV flux via IC scatterings when the gap is located well inside of the null surface. This is because the dense X-ray field will limit the particle Lorentz factors small (Paper~II), and because the less-energetic particles scatter copious IR photons into lower $\\gamma$-ray energies with large cross sections ($\\sim \\sgT$). There is room for further investigation on this issue.  \\par \\vspace{1pc}\\par  One of the authors (K. H.) wishes to express his gratitude to Drs. Y. Saito and A. K. Harding for valuable advice. He also thanks the Astronomical Data Analysis Center of National Astronomical Observatory, Japan for the use of workstations."
    },
    "0011/astro-ph0011441_arXiv.txt": {
        "abstract": "A detailed discussion of the X-ray luminosity functions of pre-main sequence and young main-sequence objects in the Taurus-Auriga-Perseus region based on the {\\em ROSAT} archive is presented. One of the main conclusions is that the XLF of classical and weak-line T Tauri stars (TTS) are different: weak-line TTS are X-ray brighter than classical TTS. Various possible biases related to the sample selection are described and ruled out as a reason for the observed discrepancies. The X-ray emission of the TTS is compared to that of the Pleiades and the Hyades to examine the influence of stellar evolution on the activity level. ",
        "introduction": "Within the last years extensive X-ray observations of star forming regions have established that the youngest among the stellar coronal X-ray sources are characterized by the strongest activity levels. A standard picture of stellar magnetic activity has developed in which the activity decreases as the stars evolve from the pre-main sequence (PMS) to the main sequence (MS), and continues to decline during the MS life.  While in some early investigations of the X-ray emission of PMS objects a power-law relation between age and stellar activity was found, other observations seem to indicate that activity remains relatively constant during the PMS and starts to drop rapidly once the MS is reached (Feigelson \\& Kriss 1989, Walter et al. 1988). Instead of being a pure age effect the decline of activity could be related to the rotational evolution of the objects: Once on the MS the stars spin down and the dynamo should become less efficient. Many observational studies concerning the connection between rotation and X-ray activity have indicated a correlation between X-ray emission and rotational velocity favoring the spin-down scenario (see e.g. Bouvier 1990; Neuh\\\"auser et al. 1995 = N95; Stauffer et al. 1997). However, these results were contradicted by other studies (Gagn\\'e et al. 1995; Micela et al. 1996). Thus, despite all efforts, the details of the relation between age, activity, and rotation are still not well understood.  An important tool to assess the strength of the X-ray activity are X-ray Luminosity Functions (XLF), i.e. cumulative distributions of X-ray luminosities $L_{\\rm x}$. Previous analyses of XLF of young stars have either focused on individual X-ray exposures of selected sky regions or were based on spatially extended but low-sensitivity observations, such as e.g. the {\\em ROSAT} All-Sky Survey (RASS).  In this contribution a systematic analysis of the XLF of young stars in Taurus-Auriga-Perseus on basis of all publicly available {\\em ROSAT} PSPC observations pointed towards that portion of the sky is presented. The stellar sample includes both PMS objects (T Tauri Stars from Taurus-Auriga), and representatives of young MS stars (the Pleiades and the Hyades cluster). A total of 106 observations are evaluated and more than 800 sources detected. The statistical analysis was performed with the ASURV package (Feigelson et al. 1985) to take proper account of upper limits for non-detections. ",
        "conclusions": "A study of the XLF from a large sample of PMS and ZAMS objects was performed on basis of pointed {\\em ROSAT} PSPC observations.  This analysis shows that the XLF are sensitive to the way in which unresolved binaries are dealt with. Two cases have been considered, in which the observed luminosity is either all attributed to the primary or distributed equally onto all components in multiples. A comparison shows that the latter assumption is consistent with the XLF of single stars. If all X-rays are assumed to come from the primary instead, the XLF of binaries are too bright.  In Taurus-Auriga the wTTS show higher $L_{\\rm x}$ than the cTTS in agreement with earlier studies in the same star forming region, but in contrast to results from other regions (e.g. Cha~I, $\\rho$~Oph). A detailed analysis was undertaken to examine possible explanations for the discrepancy of the XLF of cTTS and wTTS in Taurus-Auriga. Effects which could be ruled out are \\begin{itemize} \\item an X-ray bias in the wTTS sample, \\item different age distributions of cTTS and wTTS, \\item different spectral type distributions of cTTS and wTTS, \\item uncertainties in the X-ray emission from unresolved binaries (not discussed in the text). \\end{itemize} The only obvious parameter left with a possible influence on the X-ray emission is rotation. A comparison of slow to fast rotating wTTS in Taurus-Auriga indicates that fast rotators show higher $L_{\\rm x}$. However, the sample of stars with measured rotation periods is small (16 fast versus 12 slow rotators among the wTTS).  For a more detailed discussion of the XLF and the rotation-activity relation we refer to Stelzer \\& Neuh\\\"auser (2000)."
    },
    "0011/hep-ph0011048_arXiv.txt": {
        "abstract": "Recent models invoking extra spacelike dimensions inhabited by (bulk) neutrinos are shown to have significant cosmological effects if the size of the largest extra dimension is $R \\gtrsim 1 {\\rm \\,fm}$. We consider effects on cosmic microwave background anisotropies, big bang nucleosynthesis, deuterium and $^6$Li photoproduction, diffuse photon backgrounds, and structure formation. The resulting constraints can be stronger than either bulk graviton overproduction constraints or laboratory constraints. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011261_arXiv.txt": {
        "abstract": "We use N-body simulations combined with semi-analytic models to compute the clustering properties of modeled galaxies at $z\\sim 3$, and confront these predictions with the clustering properties of the observed population of Lyman-break galaxies (LBGs). Several scenarios for the nature of LBGs are explored, which may be broadly categorized into models in which high-redshift star formation is driven by collisional starbursts and those in which quiescent star formation dominates.  For each model, we make predictions for the LBG overdensity distribution, the variance of counts-in-cells, the correlation length, and close pair statistics.  Models which assume a one-to-one relationship between massive dark-matter halos and galaxies are disfavored by close pair statistics, as are models in which colliding halos are associated with galaxies in a simplified way. However, when modeling of gas consumption and star formation is included using a semi-analytic treatment, the quiescent and collisional starburst models predict similar clustering properties and none of these models can be ruled out based on the available clustering data.  None of the ``realistic'' models predict a strong dependence of clustering amplitude on the luminosity threshold of the sample, in apparent conflict with some observational results. ",
        "introduction": "In recent years there has been impressive growth in observations of high-redshift galaxies. The ``Lyman-break'' technique \\citep{steidel:92,madau:96,steidel:96a} makes it possible to select high-redshift candidates based on their photometric colors. Extensive spectroscopic follow-up has confirmed that this technique very reliably selects high redshift ($z \\ga 2$) galaxies \\citep{steidel:96a,steidel:96b,low:97}.  The largest sample covers the redshift range $2 \\la z \\la 3.5$, where over 1200 photometric candidates and about 900 spectra have now been obtained, mainly by Steidel and collaborators.  Similar techniques can be used to identify galaxies at even higher redshifts, although spectroscopic confirmation is more difficult.  About fifty confirmed objects exist at $4.5 \\la z \\la 5.5$ \\citep{steidel:99} and a handful at $z \\ga 5.0$ (e.g., \\citealt{weymann:98,spinrad:98}).  Our main focus in this paper will be the $z\\sim3$ LBG sample accumulated by the Steidel group, which is fairly complete to ${\\mathcal R} = 25.5$, allowing robust estimation of the clustering properties at this redshift and magnitude limit.  The correlation length of the $z\\sim3$ sample is similar to that of nearby bright galaxies ($r_0 \\sim 3$--6 \\hmpc, comoving; \\citealt{adel:98,giav:98,giav:00}; hereafter A98, G98 and G00). Within the Cold Dark Matter (CDM) hierarchical structure formation paradigm, these galaxies must therefore be much more clustered than the underlying dark-matter density field (i.e., strongly ``biased''). Moreover, because the clustering of matter increases monotonically with time, the bias of the Lyman-break galaxies must be significantly higher than that of typical galaxies at $z=0$. Although the actual level of bias and the details of its redshift dependence depend on the cosmological model and the sample selection, qualitatively this result is quite general and was pointed out by the first observational papers on LBG clustering (A98,G98) as well as numerous subsequent works. This is clearly a key property of high-redshift galaxies and must be explained by any successful theory of galaxy formation.  In the CDM framework, given a power spectrum and a cosmology, the clustering properties of dark-matter halos can be readily estimated, either by analytic methods or using N-body simulations. Numerous groups have shown, using a variety of methods, that the observed clustering and high bias of high-redshift galaxies can plausibly be reproduced in a broad range of CDM cosmologies \\citep{mofuku:96,adel:98,wech:98,jing:98,bagla:98a,baugh:98,gov:98,coles:98, moscar:98,khw:99,arnouts:99,kauf:99b,blanton:00}.  This implies that the clustering properties of LBGs are not likely to provide very discriminatory constraints on cosmology, especially as long as secure knowledge about their masses is lacking. However, there is still hope that LBG clustering may provide important constraints on galaxy formation.  A remaining central uncertainty is the association of dark-matter halos with observable galaxies.  Many previous investigations \\citep{mofuku:96,adel:98,wech:98,jing:98,bagla:98a,coles:98,moscar:98,arnouts:99} have made the simple assumption that every dark-matter halo above a given mass threshold hosts one observable LBG, and that the galaxy luminosity is closely connected with the mass of the host halo. Within the observational uncertainties at the time of publication of these earlier works, the observed number density and correlation length of the $z \\sim 3$ sample could be reproduced within this sort of scenario, provided that LBGs were associated with massive halos ($\\ga {\\rm few} 10^{11} \\hMsun$ for low-$\\Omega$ cosmologies). We shall refer to this class of models as ``Massive Halo'' models for the remainder of this paper.  \\citet[hereafter K99]{kolatt:99} investigated a very different model for LBGs, one in which \\emph{all} of the observed high-redshift galaxies are visible because they are temporarily brightened by starbursts triggered by collisions. Colliding halos were identified in a high-resolution N-body simulation, and a simple approach was used to associate these collisions with visible LBGs. Collisions between subhalos can lead to multiple LBGs within the same virialized halo. Although many of the objects in this scenario are far less massive than in the Massive Halo scenario described above, K99 showed that the correlation length of the colliding halos was comparable to that of the observed LBGs, and that the colliding halos were biased with respect to the dark matter. This demonstrated that the mass threshold of the host halos does not uniquely determine the clustering properties of a population of objects. We shall investigate a model similar to the K99 model, which we refer to as the ``Colliding Halo'' (CH) model.  Though both the Massive Halo and Colliding Halo models were able to simultaneously fit the number density and clustering properties of LBGs, they both rely on an ad hoc connection between dark-matter halos and observable galaxies, and are almost certainly too simple to be correct in detail. More detailed modeling of LBGs, relying on either semi-analytic modeling or hydrodynamic simulations to treat the physics of gas cooling and star formation, has led to a variety of different views regarding the masses and basic nature of the LBG population.  Using a semi-analytic model similar to that presented by \\citet{cafnz:94}, \\citet{baugh:98} showed that under their assumptions, LBGs are hosted by massive halos ($\\ga 10^{12} \\hMsun$), and are forming stars mainly quiescently at a moderate rate.  The correlation length of LBGs in their model was similar to that obtained in the simpler Massive Halo models and consistent with the observational estimates available at that time (see also \\citealt{gov:98}). We refer to this picture, in which LBGs are massive, quiescently star-forming objects, as the ``massive quiescent'' scenario.  Also using semi-analytic models, \\citet*[hereafter SPF]{spf:00} showed that the numbers and properties of high-redshift galaxies in such models are very sensitive to the star formation recipe adopted. They investigated three models, corresponding to three different recipes for star formation, all of which produced good agreement with local observations. In the ``Constant Efficiency Quiescent'' (CEQ) model, all star formation occurs in a quiescent mode and the star formation efficiency (i.e. the star formation rate per unit mass of cold gas) is constant with redshift.  In the ``Accelerated Quiescent'' (AQ) model, all star formation is quiescent but its efficiency scales inversely with the disk dynamical time, thus increasing rapidly at high redshift. In the third, the ``Collisional Starburst'' (CSB) model, in addition to quiescent star formation, galaxy-galaxy mergers (both major and minor) are assumed to trigger starbursts --- brief episodes in which the rate of star formation is dramatically higher than in the usual quiescent mode. The Collisional Starburst model was favored by SPF as they found that it produced the best overall agreement with the high-redshift data they investigated.  Based on hydrodynamic simulations, \\citet{khw:99} and \\citet{weinberg:00} supported a view intermediate to the massive quiescent scenario of \\citet{baugh:98} and the Collisional Starburst scenario favored by SPF, although closer to the first. \\citet{weinberg:00} found that their simulated LBGs resided within halos with a wide range of masses, but they still reproduced the strong clustering observed.  Most of the LBGs in their simulations do not appear to be undergoing starbursts, but the simulations do not have sufficient mass or spatial resolution to properly treat most of the collisions that SPF found to be important.  It is clear that regardless of whether semi-analytic or numerical techniques are used, the results of theoretical predictions about the nature of LBGs depend sensitively on the highly uncertain physics of star formation and feedback. Each of the proposed scenarios has potential problems. The simple Massive Halo models and the more detailed massive quiescent-type models seem to reproduce the observed clustering strength of LBGs, but the ``realistic'' versions of these models --- e.g., the Constant Efficiency Quiescent model of SPF --- have difficulty producing enough objects when dust extinction is included, and predict that the number density of bright galaxies should decline rapidly at higher redshift, in apparent conflict with observations (SPF).  An alternative recipe for quiescent star formation --- the Accelerated Quiescent recipe of SPF --- gives acceptable agreement with the number density of LBGs at $z \\ga 3$. However, this model has difficulty in producing enough very bright objects, and also consumes so much gas that it violates constraints from observations of Damped Lyman-$\\alpha$ systems (SPF) \\footnote[1]{The star formation recipe used in the hydrodynamic simulations is similar to the AQ model of SPF, since the gas consumption timescale scales with the local dynamical time.  The mass resolution is not good enough to tell conclusively whether the large amount of high-redshift gas consumption results in the same problem with matching the DLAS abundance, but \\citet{gardner:99} argue based on an analytic extention of the mass resolution that this is not a serious problem.}.  In addition, because LBGs are found in smaller mass halos, it was not clear whether they would be clustered enough to match the data.  The Colliding Halo model of K99 was shown to reproduce the clustering of LBGs on scales of several Mpc, but it may be {\\em too} clustered on smaller scales, and thus overpredict the number of close pairs \\citep{mmw:99}. The clustering properties of the more detailed Collisional Starburst model of SPF have not been checked until the present work, but could suffer the same problem.  Also, there is a suggestion that the clustering strength of observed LBGs depends on the luminosity threshold of the sample (\\citealt{steidel:98rs}, hereafter S98; G00), with brighter galaxies being more strongly clustered. Because of the expected large dispersion in the relationship between mass and luminosity, starburst models might have difficulty producing a strong trend of this sort.  The goal of this paper is to test a set of models covering the full range of previously proposed scenarios for the nature of LBGs, from the very simple Massive Halo and Colliding Halo models to the more ``realistic'' models mentioned above, and to determine which of them, if any, can be ruled out by comparing their predicted clustering properties with the available data at $z\\sim 3$. The number of observable objects per halo (the occupation function) is calculated using both simple analytic prescriptions and results from the semi-analytic models of SPF.  Large-volume dissipationless N-body simulations are used to calculate the expected clustering properties of halos at $z=3$, and the calculated occupation functions are used to convert this into predictions for the clustering properties of observable galaxies (this is similar to the approach used by \\citealt{kns:97} and \\citealt{benson:00}). We mimic the observational selection effects as closely as possible, apply them to model galaxies, and then compare these predictions to the data in the ``observational plane''.  Throughout the paper, we focus on one cosmology, the currently popular \\LCDM\\ model, with matter density $\\Omega_m = 0.3$, vacuum energy density $\\Omega_{\\Lambda} = 0.7$, and a Hubble parameter $h=0.7$, where $H_0 = 100h \\kmsmpc$.  The outline of the paper is as follows.  We begin (\\S\\ref{sec:mw-lc}) with an analytic investigation of clustering for two extreme models for LBGs: Massive Halos and Colliding Halos. In \\S\\ref{sec:data}, we discuss the data that we will use for comparison. In \\S \\ref{sec:sims} we discuss the N-body simulations we use to derive halo clustering properties, and the models that are used to populate these halos with galaxies.  In \\S \\ref{sec:results}, we present the statistics used and the results of our comparison.  In \\S \\ref{sec:compare}, we compare the Colliding Halo model and the more detailed semi-analytic Collisional Starburst model and determine which elements of the models are responsible for differences in their behavior. We discuss our results and conclude in \\S \\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu}  We have investigated a range of models, spanning the most extreme versions of previously proposed scenarios for Lyman-break galaxies, and compared the predicted clustering properties of galaxies at $z=3$ with those of Lyman-break galaxies with spectroscopic redshifts from the ground-based sample of Steidel, Adelberger, and collaborators. We investigated two simple models for assigning observable galaxies to halos, the Massive Halo model, which associates one LBGs with each massive halo, and the Colliding Halo model, which associates an LBG with each halo collision. In addition, we investigated three ``realistic'' models based on the semi-analytic models of SPF, which are differentiated by their assumptions about the efficiency of transforming cold gas into stars.  All five models are normalized to produce the observed, incompleteness corrected number density of LBGs \\citep{steidel:99}.  For each model, we then compute the halo occupation function, or the number of observable galaxies per halo as a function of the host halo mass. These occupation functions are then used to populate dark-matter halos in a large, dissipationless \\LCDM\\ N-body simulation with galaxies. We select galaxies, attempting to mimic the observational selection techniques, and then use the resulting catalogs to calculate the clustering properties of LBGs in each scenario.  Clustering statistics that smooth over a relatively large region (such as the overdensity distribution in cells of $\\sim 12$ \\hmpc\\ on a side, or the correlation length) are fairly non-discriminatory for all five models.  No models can be strongly ruled out based on the correlation length; the most clustered is the Colliding Halo model, which is about $2\\sigma$ away from the data.  The correlation lengths for the rest of the models are in reasonable agreement with the data, when compared using the same technique.  However, there is still uncertainty in the determination of the true correlation length of LBGs; if it is really closer to $3$ \\hmpc\\, rather than 4--5 \\hmpc, then this is probably an indication that LBGs must be hosted by smaller mass halos than any of our models predict.  A change of a factor of two in the true number density of the objects seen as LBGs could possibly account for this. Alternatively, for the ``realistic'' models, the uncertainties in the modeling of star formation, feedback, and dust extinction, the stellar initial mass function, etc., are likely to be sufficient to allow for some adjustments that would change the occupation function in just the right way to satisfy all the constraints investigated here.  One possibility, which we have not investigated in detail here, is that the most rapidly star-forming galaxies suffer so much dust extinction that they are not included in the observational sample.  In this case, the proportion of visible galaxies in massive halos would be further reduced --- thus lowering the correlation length.  The hope, of course, is that the empirical understanding gained by this sort of schematic modeling will lead to constraints that will eventually bring about progress towards a deeper understanding of these physical processes.  The fraction of galaxies in close pairs, which is a good measure of the number of multiple galaxies in a single halo, proved to be a better discriminator between models than the statistics that measure clustering on larger scales. This is because this quantity is quite sensitive to the slope $S$ of the halo occupation function at the high mass end, $N_g \\propto M^S$.  This statistic can discriminate between the most extreme models, the simple Massive Halo model ($S=0$), which underproduces close pairs by more than $1\\sigma$, and the Colliding Halo model ($S=1.1$), which overproduces them by several $\\sigma$ (see Figure \\ref{fig:pairs}).  The Collisional Starburst model, which has a halo occupation slope of $S\\sim 0.7$ produces good agreement with the close pair fraction of the observations; the other two ``realistic'' semi-analytic models yield a slightly higher slope of $S\\sim 0.8$, and a higher fraction of pairs than the data, but are not strongly ruled out by the current data.  We also examined the dependence of clustering on luminosity. There have been suggestions that the correlation length derived from the observations depends on the magnitude limit of the sample, with brighter galaxies being more strongly clustered (S98, G00). This has been interpreted as evidence for a tight connection between halo mass and galaxy luminosity, which would seem to disfavor models with stochastic star formation such as collisional starbursts. However, when all available measurements of the correlation length from the literature are compiled, it seems clear that the strength of this observational trend depends greatly on which subset of the observations one chooses to include.  Until a single high-redshift sample exists that is large enough in volume to measure the clustering of bright galaxies without being dominated by shot-noise, and deep enough to simultaneously measure the clustering of much fainter galaxies, it will be difficult to place strong limits on the models.  In the meantime, for the models, we find that \\emph{both} the quiescent and collisional starburst models display only a weak trend of correlation length with luminosity, much weaker than that suggested by the observational analysis of G00, or predicted by the Massive Halo model. As expected, the relation between galaxy luminosity and halo mass has much larger scatter in starburst models than in quiescent models.  However, since massive halos can have a number of subhalos, with a variety of masses, both models actually have quite a large scatter in the luminosities of galaxies that reside in these host halos, which are what determine the clustering properties.  Any hierarchical clustering model will have multiple subhalos within the massive halos, with a range of masses, several of which are likely to correspond to observable galaxies. In fact, the observed fraction of galaxies in close pairs seems to indicate that many of the halos do host more than one visible galaxy, since the number of pairs in the Massive Halo model is more than $1\\sigma$ lower than the data (the first bin in Figure~\\ref{fig:pairs}). It is a concern that there is some disagreement between the subhalo multiplicity function of the semi-analytic models compared to that in the simulations; if in fact the number of subhalos predicted by the semi-analytic treatment is too high, the dependence of clustering on luminosity would be somewhat stronger.  We found that two versions of the collisional starburst scenario gave different results for the halo occupation function and tried to understand why. The Colliding Halo model, based on properties of collisions identified in high-resolution N-body simulations, had many more LBGs in massive halos than the semi-analytic model, and a steeper halo occupation function.  Although the difference could have been caused by an inaccurate treatment of halo mergers in the semi-analytic analysis, we found that this was not the case.  Instead, the discrepancy between this model and the semi-analytic Collisional Starburst model arises because of the detailed modeling of gas processes and starburst efficiency in the latter.  Namely, high-mass halos in the semi-analytic model are less efficient at producing bright galaxies.  This is primarily due to two effects: massive halos collapse more recently and have not had time to cool as large a fraction of their gas, and (the more dominant effect) bursts in minor mergers are suppressed when the primary galaxy already contains a prominent bulge. This recipe was adopted by SPF to attempt to capture the behavior found by \\citet{mihos:96}, based on hydrodynamic simulations of merging galaxies. Apparently, a morphology-density relation is already in place in our models at $z \\sim 3$, and this has a significant impact on the predicted properties of observable galaxies. It should be noted that although SPF based their recipe for burst efficiency on the best simulations that were available at the time, the sensitivity of our results to this ingredient is a concern. A more extensive set of simulations, with initial conditions chosen to better represent $z \\sim 3$ galaxies and covering a larger region of parameter space, is in progress \\citep{somerville:00} and will hopefully improve our understanding of this process.  In the meantime, one should be aware that the predictions for clustering in this model are especially sensitive to the assumed efficiency of converting gas into stars during a starburst.  In summary, although one might have expected the clustering properties of galaxies on intermediate scales (i.e., the correlation length) to provide a strong discriminator between models of galaxy formation, in fact we find that even extreme models yield similar results for this statistic --- although models with a very steep occupation function are only marginally acceptable.  Clustering statistics that probe smaller scales are a better way to discriminate between models whose halo occupation slopes are different, but even with this statistic, none of the more realistic models can be strongly ruled out with the data sample used here.  Interestingly, we found that the halo occupation slope was shallower in the Collisional Starburst model than in the massive quiescent models investigated, laying to rest concerns that this model would comparatively have too many close pairs.  When combined with the failings of the other two (quiescent) models discussed in SPF, it still appears that the Collisional Starburst provides the best agreement with all the available data.  Still, we emphasize that many ingredients of the modeling of collisional starbursts remain highly uncertain, and if indeed this process is responsible for producing most of the galaxies observed at high redshift, further investigation will be crucial.  It should also be emphasized that although the current observational situation is still too uncertain to unambiguously determine how these high-redshift galaxies are related to the underlying dark halos, there is hope that in the future, a combination of observations on number density, and both small-scale and large-scale clustering, should be able to determine the halo occupation function, which we then can hope to explain with physical models for galaxy formation."
    },
    "0011/astro-ph0011057_arXiv.txt": {
        "abstract": "A sudden steepening of the cosmic-ray energy spectrum ( the knee ) is observed at an energy of about 3 PeV (1 PeV = 10$^{15}$eV). The experimental study of the PeV cosmic rays has intensified greatly during the last 3 years. The recent results on extensive air showers allow us to conclude that: a) the knee has an astrophysical origin; b) the `sharpness' and the fine structure of the knee rule out `Galactic Modulation' as the origin of the knee; c) most likely the knee is the result of the explosion of a single, recent, nearby supernova. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011327_arXiv.txt": {
        "abstract": "On 30.14 March 1997 we observed the EUV spectrum of the bright comet C/1995 O1 (Hale-Bopp) at the time of its perihelion, using our EUVS sounding rocket telescope/spectrometer.  The spectra reveal the presence H Ly $\\beta$, O$^+$, and, most notably, Argon.  Modelling of the retrieved Ar production rates indicates that comet Hale-Bopp is enriched in Ar relative to cosmogonic expectations.  This in turn indicates that Hale-Bopp's deep interior has never been exposed to the 35--40 K temperatures necessary to deplete the comet's primordial argon supply. ",
        "introduction": "As well preserved, ancient relics of the chemistry present at the formation stage of the outer solar system, comets provide one of the most valuable tools available for understanding the formation processes and conditions extant when the planets formed, $\\approx$4.5 Gyr ago \\citep{Mea93}.  Among the most long-standing observational goals for understanding cometary origins has been the search for cometary noble gases.  Owing to their combination of high volatility and disaffinity to chemical reactions, noble gases provide a key diagnostic to the thermal history of cometary ices.  More specifically, the members of the He, Ne, Ar, Kr sequence display successively higher sublimation temperatures; as such, they provide a series of thermometers that can be exploited to constrain the thermal history, and therefore the sites of cometary origins.  Although the interpretation of noble gas abundances in cometary comae is complicated by the details of their trapping and release efficiencies in cometary ice \\citep{Oea91}, their detection has nonetheless been highly desired. Unfortunately, however, although significant upper limits revealing He depletions of $\\sim$10$^4$ in comet Austin \\citep{Sea92} and Ne depletions of $\\sim$25 in comet Hale-Bopp \\citep{Kea97} have been obtained, no detection of any cometary noble gas emmanating from within a comet has previously been obtained.  He, Ne, Ar, and Kr each have resonance transitions in the far and extreme ultraviolet. Among the noble gases, argon offers a particularly good combination of comparatively high cosmogonic abundance, moderate sublimation temperature, and good UV resonance fluorescence efficiency; together these properties suggested some time ago \\citep{Sea92, Mea93} that argon may be the easiest noble gas to detect in comets.  We took advantage of the apparition of the unusually active and bright comet C/1995 O1 (Hale-Bopp) in order to conduct a new and more sensitive search for argon in comets. ",
        "conclusions": ""
    },
    "0011/astro-ph0011111_arXiv.txt": {
        "abstract": "We present disk wind model calculations for the broad emission lines seen in the ultraviolet spectra of the X-ray binary Hercules X-1. Recent HST/STIS observations of these lines suggest that they are kinematically linked to the orbital motion of the neutron star and exhibit a red-shifted to blue-shifted evolution of the line shape during the progression of the eclipse from ingress to egress which is indicative of disk emission.  Furthermore, these lines are single-peaked which implies that they may be formed in a disk wind similar to those we have proposed as producing the broad emission lines seen in the UV spectra of active galactic nuclei.  We compute line profiles as a function of eclipse phase and compare them to the observed line profiles.  Various effects may modify the appearance of the lines including resonant scattering in the wind itself, self-shadowing of the warped disk from the central continuum, and self-obscuration of parts of the disk along the observer's line-of-sight.  These latter two effects can cause orbital and precessional phase dependent variations in the emission lines.  Hence, examination of the line profiles as a function of these phases can, in principle, provide additional information on the characteristics of the disk warp. ",
        "introduction": "Hercules X-1 is one of the best-studied objects in the X-ray sky and displays a wealth of behavior which make it a unique laboratory for investigating the physics of accretion disks in binary systems.  It emits X-ray pulsations with a 1.24\\,s period indicating that the compact object is a rapidly rotating neutron star whose magnetic field is modulating the X-ray emission (Tananbaum et al.\\ 1972).  In addition, its hard X-ray spectrum shows features consistent with resonant cyclotron absorption in a strong magnetic field, $B \\ga 10^{12}\\,$G, as would be expected at the neutron star surface (Tr\\\"umper et al.\\ 1978; Soong et al.\\ 1990).  The system has a 1.7\\,day orbital period, and the associated X-ray eclipses imply a highly inclined ($i \\ga 80^\\circ$) circular orbit and an intermediate-sized companion, HZ Herculis, with mass $M \\simeq 2.2\\,M_\\odot$ (Deeter et al.\\ 1981).  The X-ray light curve of Her X-1 is also modulated on a 35\\,d period (Giacconi et al.\\ 1973) and consists of a ``main-high'' state lasting 10\\,d, followed by a 10\\,d quiescent period, a subsequent ``short-high'' state of 7--8\\,d, and finally another quiescent period of 7--8\\,d.  In addition, the X-ray light curve shows pre-eclipse dips which occur at the beat frequency of the orbital and 35\\,d periods. The 35\\,d cycle of Her X-1 and the behavior of the pre-eclipse dips have been studied in great detail by Scott \\& Leahy (1999) using RXTE/ASM data.  The 35\\,d modulation of the X-ray flux has been attributed to a precessing accretion disk which is tilted and/or warped so that the disk itself periodically obscures the central X-ray source (Katz 1973).  The origin of such a warp is not clear, but relatively recent theoretical work has shown that both the warp and the disk precession may be due to a radiation-driven warping instability (Pringle 1996; Maloney, Begelman, \\& Pringle 1996).  In these models, radiation from the central source illuminates the disk and is absorbed.  Since this central illumination is radially directed, it contributes no net torque to any given annulus of the disk.  However, if the disk annulus is optically thick, then the net pressure for re-radiated emission will be normal to the surface of the disk.  Therefore, if the disk annulus is illuminated non-axisymmetrically, a net torque can force the annulus out of the nominal disk plane.  Non-axisymmetric illumination may arise if the local angle of tilt of the disk varies with position, i.e., if the disk is warped.  This is the basis for the radiation-driven instability.  Alternatively, such torques may be produced in a similar fashion by non-axisymmetric winds driven from the disk surface by the X-ray heating (e.g., Schandl \\& Meyer 1994; Schandl 1996).  However, the nature of the underlying instability is essentially the same.  Radiation-driven warping is an appealing mechanism since it allows for global precessional modes which appear to be required by the regularity of the long term periodicity, and it also provides a natural time scale for the precession (Maloney, Begelman, \\& Pringle 1996).  For Her X-1, a plausible estimate for this time scale can easily account for the observed period: $\\tau_{\\rm prec} \\sim 30\\,(\\alpha/0.1)^{-1}(\\epsilon/0.1)^{-3}$ days, where $\\alpha$ is the usual disk viscosity parameter and $\\epsilon = L/\\dot{M}c^2$ is the efficiency of the accretion (Maloney \\& Begelman 1997).  The shape of the X-ray light curve over the 35\\,d period can provide some constraints on the characteristics of the disk warp and hence on the physical mechanisms which produce it.  Schandl \\& Meyer (1994) and Schandl (1996) performed dynamical calculations of a disk warp driven by coronal winds from the disk surface which are powered by heating due to the central X-rays.  This model accounts for the presence of the two high states via a combination of obscuration by the disk itself and by an optically thick corona overlying it.  These authors did not consider the effect of re-radiation torques in their calculations, which may be at least as important as the torques due to a wind (see discussion in Maloney, Begelman, \\& Pringle 1996).  In contrast, Scott, Leahy, \\& Wilson (2000) have recently described a completely phenomenological model for the tilt and phase of the inner and outer disk annuli which are adjusted to fit the main-high and short-high intervals.  In addition to differing from the Schandl-Meyer geometry, their model accounts for the X-ray spectral softening seen during both high state turn-ons as being due to absorption by cold material in the outer disk.  Since the Schandl-Meyer model attributes the short-high state to emergence of the X-ray source from behind the hot, optically thick disk corona, it cannot account for this spectral variation.  Although the X-ray observations can clearly place some constraints on the nature of the putative disk warp, additional observational diagnostics would be helpful.  In addition to driving warps, X-ray illumination of the disk can also produce emission line radiation.  Several authors have performed calculations of the disk structure and radiative transfer for an X-ray illuminated disk (Raymond 1993; Ko \\& Kallman 1994).  These models predict the existence of an accretion disk corona as well as a chromospheric layer which produces emission from a variety of lines ranging from the Fe K lines in the X-rays to UV lines such as \\Civ, \\Nv, \\Siiv, \\Heii, and \\Ov.  Although the predicted line ratios appear to agree substantially with those observed in well-studied objects such as Her X-1 and Sco X-1, since they are produced in a disk, the lines are expected to be double-peaked, contrary to observations (Vrtilek et al.\\ 1991; Boroson et al.\\ 1996, 2000a).  We present a model for the broad single-peaked UV emission lines in which they are formed in an accretion disk wind.  The underlying mechanism is the same as the one we have proposed for the broad lines in active galactic nuclei (Chiang \\& Murray 1996; Murray \\& Chiang 1997) and the single-peaked emission lines in cataclysmic variables (Murray \\& Chiang 1996): Large radial shears, due to the wind, in line emitting material at the near and far sides of the accretion disk allow line photons with small Doppler shifts to escape and produce single-peaked emission lines.  The remainder of this paper is organized as follows: In \\S~2, we review the recent observations of the UV emission lines in Her X-1, and in particular, discuss the evidence for a disk origin for the broad UV lines.  In \\S~3, we review the mechanism for producing broad single-peaked emission lines in disk winds.  In \\S~4 we apply these calculations to fit the July 1998 HST/STIS data obtained during the eclipse phases of Her X-1.  We also show how disk warping can affect the appearance of the lines as a function of orbital and precessional phase.  Finally, in \\S~5, we present our conclusions. ",
        "conclusions": "We have presented calculations of the broad UV emission line profiles which are formed in an accretion disk wind for the X-ray binary Her X-1.  Calculations performed assuming a flat accretion disk reproduce the eclipse light curve shapes for the July 1998 HST/STIS data reasonably well, although previous analysis of these data revealed out-of-eclipse variability in the broad line fluxes which are probably due to time-dependent variations in the emitting properties of the disk material.  The flat disk line profiles also do a reasonable job of reproducing the shapes of the emission lines during eclipse as they evolve from red-shifted on ingress to blue-shifted on egress.  The flat disk model does not provide a perfect description of the line shapes however.  At orbital phases $\\phiorb = 0.902$--0.916, the observed profiles are more red-shifted than the model calculations. This may be due to scattering in an outflowing wind as evidenced by the residual P\\,Cygni profiles for these phases (Fig.~\\ref{early_phase_residuals}).  This explanation is attractive since P\\,Cygni line profiles appear in the 1999 July HST/STIS spectra of Her X-1 at orbital phases $\\phiorb \\sim 0.5$ (Boroson, Kallman, \\& Vrtilek 2000b).  Alternatively, the relative red-shifts of the observed line shapes may be due to self-obscuration of the blue-shifted portions of a warped disk (Figs.~\\ref{SM_warp_line_profiles} and~\\ref{SLW_warp_line_profiles}). Such self-obscuration may also be able to account for the line shape differences which are seen in the high resolution spectra at $\\phiorb = 0.057$ which the flat disk model consistently overpredicts (Figs.~\\ref{HeII_profiles}, \\ref{Civ_profiles}, \\ref{Nv_profiles}). However, B00 note that there is evidence for narrow absorption in the \\Nv\\ and \\Civ\\ profiles at a Doppler shift of $-500$\\kms\\ which may significantly modify those line profiles at $\\phiorb = 0.057$ since the overall Doppler shifts of those lines are also $\\sim -500$\\kms\\ at this phase.  This absorption also appears at phases $\\phiorb = 0.092$, 0.132, and 0.171 at the same velocity, so it is unlikely to be associated with the disk either as an obscuring warp or outflow, as the overall Doppler shift would then vary by $\\sim 100$\\kms\\ over these phases (B00).  The main difficulty which this model faces is how a wind with the required characteristics can be launched from the surface of the disk. In previous analytic and numerical calculations of disk winds in AGNs and CVs, line driving provided the mechanism for accelerating the wind in which most of the initial acceleration was due to radiation produced in the disk itself (e.g., Murray et al.\\ 1995; Proga, Stone, \\& Drew 1998).  Although our treatment does not specifically require radiative acceleration to play a substantial role in producing the wind, large poloidal velocity gradients in the emission line region are a necessary ingredient.  In the case of X-ray binaries such as Her X-1, it has been posited that any wind emerging from the disk would likely be thermally driven and have its origins in a Compton-heated corona which has been irradiated by the central X-ray source (Begelman, McKee, \\& Shields 1983).  The scale height for such a wind should be comparable to the disk radius.  This violates the underlying assumption in these calculations of a geometrically thin, albeit warped, disk.  Furthermore, given these large vertical scale heights, it is unlikely that there will be poloidal velocity gradients which are comparable to the Keplerian velocity gradients.  An alternative mechanism for driving a wind from a disk is via magneto-centrifugal acceleration (e.g., Blandford \\& Payne 1982).  In this model, the magnetic field lines thread the accretion disk and are anchored to the disk by flux freezing.  If the field lines rotate rigidly with the disk and make an angle greater than $60^\\circ$ with the disk normal, then material will be centrifugally accelerated along the field lines.  If the magnetic field is largely poloidal then the acceleration of the disk material can be estimated by the condition that the angular velocity is constant along the field line.  This leads to radial shears which are comparable to the Keplerian shears in the disk and single-peaked emission lines can be produced (Chiang \\& Murray 1997).  However, in the self-similar solution of Blandford \\& Payne (1982; see also Emmering, Blandford, \\& Shlosman 1992), the magnetic fields near the disk are not mostly poloidal since the inertia of the accelerated material actually distorts the field lines and causes them to have a significant toroidal component in the $-\\hat{\\phi}$-direction.  The azimuthal shears are actually enhanced; the radial shears are much smaller than they would be if the field lines were poloidal; and as a result, the emission line profiles are strongly double-peaked (Chiang 2000).  It is far more likely in the case of Her X-1 that radiative acceleration does play a role in the formation of the broad emission lines.  The existence of a chromospheric layer in which atoms with strong resonance lines are abundant implies that line driving pressure may well be important in determining the vertical structure of a X-ray irradiated disk atmosphere.  For \\Ostars, Castor, Abbott, \\& Klein (1975; hereafter CAK) argued that even in a static atmosphere the force due to lines will greatly exceed that due to gravity, making a line-driven wind inevitable.  A cursory examination of the properties of the X-ray irradiated disk atmospheres computed by Raymond (1993) and Ko \\& Kallman (1994) reveals that similar conditions exist which may allow significant line-driving to occur.  Following CAK, we use a force multiplier formalism to describe the effects of line driving. Here the enhancement factor relative to electron scattering, the force multipler ${\\cal M}$, is modeled as an inverse power of the electron scattering optical depth: ${\\cal M} = k t^{-\\alpha}$, where $k$ and $\\alpha$ are parameters which depend on the driving continuum and on the composition and ionization state of the irradiated gas.  For a static atmosphere, the electron scattering optical depth is $t \\simeq \\sigma_e\\rho l$ where $\\sigma_e$ is the electron scattering opacity, $\\rho$ is the mass density, and $l$ is the appropriate length scale which, for a disk wind, is $l_{\\rm em}$. In the case of a X-ray irradiated disk atmosphere, we can estimate the relevant value of $t$ from the line optical depths quoted by Raymond and Ko \\& Kallman.  The latter authors find line optical depths of $\\ga 10^4$.  For resonance lines such as \\Civ\\ this implies an electron scattering depth of $t \\sim 10^{-5}$ for the chromospheric layer assuming a gas temperature of $10^4$\\,K.  For \\Ostars, such a small value of $t$ corresponds to the regime where ${\\cal M}$ has saturated at ${\\cal M}_{\\rm max} \\sim 10^3$ essentially because of the oscillator strength sum rule (CAK; Abbott 1982).  At disk radii $r \\sim 10^{10}$\\,cm, Ko \\& Kallman (1994) find disk photosphere temperatures in the range $T_{\\rm disk} =10^4$--$10^5$\\,K, which is comparable to that of \\Ostars\\ making our use of this saturation value reasonable.  Using this value of the force multiplier and $F = \\sigma T_{\\rm disk}^4$ for the driving flux, the acceleration due to radiation pressure from the disk is $f_{\\rm rad} = \\sigma_e F {\\cal M}/c \\simeq 6\\times 10^5 (T_{\\rm disk}/3\\times10^4{\\rm K})^4$\\,cm\\,s$^{-2}$.  This exceeds the vertical component of gravity $f_g \\simeq 10^5$\\,cm\\,s$^{-2}$ assuming a central mass of 1.3\\,$M_\\odot$, a disk radius of $10^{10}$\\,cm and a vertical height for the chromosphere of $10^9$\\,cm.  These arguments suggest that it would be worthwhile to perform calculations of X-ray irradiated disk atmospheres which include the effects of radiative acceleration.  Such calculations would likely result in substantial changes to the disk vertical structure, line optical depths, and local energy balance, and may be able to account for the appearance of marginally optically thick emission lines such as those we have inferred for \\Nv.  Furthermore, if radiatively driven winds do result, they may be important for producing a disk warp.  For a given incident flux of radiation, winds can be more efficient at driving warps than radiation pressure (Pringle 1996).  However, given the large scale heights of thermally driven winds such as those examined by Schandl \\& Meyer, it is difficult to understand the manner in which the wind back pressure is communicated down to the disk plane in order to produce non-axisymmetric torques on the underlying disk annuli.  By contrast, the vertical length scale of a radiatively driven wind will be much smaller than that of a thermally driven wind since the gas does not have to reach the escape temperature in order to become unbound.  So it seems plausible that radiative acceleration plays a role in the dynamics of the line emitting regions of the disks in X-ray binaries. The appearance of broad single-peaked lines in the UV spectra of Her X-1 and the abundant evidence from the eclipse light curves and line profiles implying a kinematic connection with an accretion disk makes this scenario even more attractive."
    },
    "0011/astro-ph0011237_arXiv.txt": {
        "abstract": "Infrared (IR) dust emission from galaxies is frequently used as an indicator of star formation rate (SFR). However, the effect of the dust-to-gas ratio (i.e., amount of the dust) on the conversion law from IR luminosity to SFR has not so far been considered. Then, in this paper, we present a convenient analytical formula including this effect. In order to obtain the dependence on the dust-to-gas ratio, we extend the formula derived in our previous paper, in which a theoretical formula converting IR luminosity to SFR was derived. That formula was expressed as ${\\rm SFR}/(M_\\odot~{\\rm yr}^{-1})=\\{ 3.3\\times 10^{-10}(1- \\eta )/(0.4-0.2f+0.6\\epsilon )\\} (L_{\\rm IR}/L_\\odot )$, where $f$ is the fraction of ionizing photons absorbed by hydrogen, $\\epsilon$ is the efficiency of dust absorption for nonionizing photons, $\\eta$ is the cirrus fraction of observed dust luminosity, and $L_{\\rm IR}$ is the observed luminosity of dust emission in the 8--1000-$\\mu$m range. Our formula explains the IR excess of the Galaxy and the Large Magellanic Cloud. In the current paper, especially, we present the metallicity dependence of our conversion law between SFR and $L_{\\rm IR}$. This is possible since both $f$ and $\\epsilon$ can be estimated via the dust-to-gas ratio, which is related to metallicity. We have confirmed that the relation between the metallicity and the dust-to-gas ratio is applied to both giant and dwarf galaxies. Finally, we apply the result to the cosmic star formation history. We find that the comoving SFR at $z\\sim 3$ calculated from previous empirical formulae is underestimated by a factor of 4--5. ",
        "introduction": "When we want to know the entire evolutionary history of a galaxy, its color and metallicity are important diagnostic quantities (e.g., Tinsley \\cite{tinsley72}). The evolution of the color and the metallicity results from the superposition of successive star formation during the lifetime of the galaxy. Thus, if we wish to reveal the evolution of galaxies, we always need to estimate the star formation rate (SFR) on a galaxy-wide scale. There are many methods to estimate the SFR of galaxies from observational quantities (Kennicutt \\cite{kennicutt98a}). In this paper, we are  especially interested in the conversion formula from infrared (IR) luminosity of the galaxies to their present SFR. The IR luminosity originates from dust that is heated by stellar radiation field. Kennicutt (\\cite{kennicutt98b}) estimated SFR from IR luminosity for starburst galaxies. Buat \\& Xu (\\cite{buat96}) adopted an empirical approach by utilizing the observed relation between the ultraviolet (UV) luminosity and the IR luminosity.  Recently, Inoue et al.\\ (\\cite{inoue00a}, hereafter IHK00) have derived a theoretical conversion formula from IR luminosity to the SFR by developing a standard model of H {\\sc ii} regions by Petrosian et al.\\ (\\cite{petrosian72}). The present paper examines how metallicity (or dust-to-gas ratio) affects the conversion factor of IHK00. It is possible because the effect of the metallicity is well parameterized in their formula. Here, we present it as \\begin{eqnarray} \\frac{{\\rm SFR}}{M_\\odot~{\\rm yr}^{-1}}= \\frac{3.3\\times 10^{-10}(1-\\eta )}{0.4-0.2f+0.6\\epsilon}\\, \\frac{L_{\\rm IR}}{L_\\odot}\\, ,\\label{inoue} \\end{eqnarray} where $f$ is the fraction of ionizing photons absorbed by hydrogen, $\\epsilon$ is the efficiency of dust absorption for nonionizing photons, $L_{\\rm IR}$ is the observed luminosity of dust in the wavelength range of 8--1000 $\\mu$m, and $\\eta$ is the cirrus fraction of $L_{\\rm IR}$. In Eq.\\  (\\ref{inoue}), $f$ and $\\epsilon$ depend on the dust-to-gas ratio\\footnote{In this paper, we focus our main attention on only $f$ and $\\epsilon$, although we should keep in mind that the ($1-\\eta$) may depend on the dust-to-gas ratio. The effect of ($1-\\eta$) on SFR is discussed in Appendix A.}. According to Hirashita (1999a, b), the dust-to-gas ratio is expressed as a function of metallicity (see also Dwek \\cite{dwek98}; Lisenfeld \\& Ferrara \\cite{lisenfeld98}). Therefore, we find easily that the relation between $L_{\\rm IR}$ and SFR depends on metallicity via $f$ and $\\epsilon$.  In this paper, we quantitatively examine the importance of the metallicity to our conversion law, the conclusion being that the metallicity dependence is not always negligible when we compare SFRs of very young galaxies and those of present galaxies. First, we examine the formula by IHK00 in Sect.\\ \\ref{sec:sfr}, focusing on the effect dust-to-gas ratio. Then, in Sect.\\ \\ref{sec:metal}, we present the metallicity dependence of the IHK00's formula. In Sect.\\ \\ref{sec:cosmic}, we extend our discussion to the cosmic star formation history. Finally, we summarize our conclusions in Sect.\\ \\ref{sec:conclusion} . ",
        "conclusions": "\\label{sec:conclusion}  Based on IHK00's formulation, we consider the factor $C_{\\rm IR}$ in the conversion formula between the IR luminosity and the SFR (Eqs.\\ (\\ref{eq:conversion}) and (\\ref{eq:c_ir})). The factor $C_{\\rm IR}$ becomes $1.7\\times 10^{-10}$, $5.6\\times 10^{-10}$, and $7.9\\times 10^{-10}$ $M_\\odot~{\\rm yr}^{-1}~L_\\odot^{-1}$ for the metallicity of 1, 0.1, 0.01 times the solar value, respectively. Thus, $C_{\\rm IR}$ differs by a factor of 4--5 in the range. Importantly, applicability of our formula is observationally confirmed by IREs of the Galaxy and the LMC. Applying our result to the cosmic star formation history, we have found that $C_{\\rm IR}$ may be about 4--5 times larger at $z\\sim 3$ than that at $z\\sim 0$. This factor of 4--5 is larger than the uncertainty caused by the initial mass function of stars (a factor of 2). Thus, we should carefully consider the chemical evolution of galaxies if we determine the cosmic star formation history from the dust emission within a factor of 4--5."
    },
    "0011/astro-ph0011001_arXiv.txt": {
        "abstract": "We study the angular-momentum profiles of a statistical sample of halos drawn from a high-resolution $N$-body simulation of the $\\Lambda$CDM cosmology. We find that the cumulative mass distribution of specific angular momentum $j$ in a halo of mass $M_{\\rm v}$ is well fit by a {\\it universal~} function, $M(<j) = M_{\\rm v} \\mu j/(j_0+j)$.  This profile is defined by one shape parameter ($\\mu$ {\\em or} $j_0$) in addition to the global spin parameter $\\lambda$.  It follows a power-law $M(<j) \\propto j$ over most of the mass, and flattens at large $j$, with the flattening more pronounced for small values of $\\mu$ (or large $j_0$ at a fixed $\\lambda$).  Compared to a uniform sphere in solid-body rotation, most halos have a higher fraction of their mass in the low- and high-$j$ tails of the distribution.  High-$\\lambda$ halos tend to have high $\\mu$ values, corresponding to a narrower, more uniform $j$ distribution. The spatial distribution of angular momentum in halos tends to be cylindrical and is well-aligned within each halo for $\\sim 80\\%$ of the halos.  The more misaligned halos tend to have low-$\\mu$ values.  When averaged over spherical shells encompassing mass $M$, the halo $j$ profiles are fit by $j(M) \\propto M^s$ with $s=1.3 \\pm 0.3$. We investigate two ideas for the origin of this profile. The first is based on a revised version of linear tidal-torque theory combined with extended Press-Schechter mass accretion, and the second focuses on $j$ transport in minor mergers.  Finally, we briefly explore implications of the $M(<j)$ profile on the formation of galactic disks assuming that $j$ is conserved during an adiabatic baryonic infall. The implied gas density profile deviates from an exponential disk, with a higher density at small radii and a tail extending to large radii. The steep central density profiles may imply disk scale lengths that are smaller than observed.  This is reminiscent of the ``angular-momentum problem\" seen in hydrodynamic simulations, even though we have assumed perfect $j$ conservation.  A possible solution is to associate the central excesses with bulge components and the outer regions with extended gaseous disks. ",
        "introduction": "\\label{sec:intro}  \\label{sec:intro} \\footnotetext[0]{Hubble Fellow}  The origin of the distribution of mass and angular momentum in disk galaxies is still an open issue, despite its long history.  The archetypical model of Eggen, Lynden-Bell \\& Sandage (1962) provides a useful framework, in which a disk galaxy (the Milky Way) forms by centrifugal support of a collapsing gas cloud. Mestel (1963) added the assumption that the specific angular momentum of each mass element, $j$, is conserved during the collapse, and demonstrated that the final mass profile of the disk could then be related to the initial mass and angular momentum distribution of the gas. The key is that the cumulative mass with specific angular momentum less than $j$, $M(<j)$, is preserved during the collapse. Crampin \\& Hoyle (1964) (also Innanen 1966; Freedman 1970) then realized that the observed exponential density profiles of disk galaxies in circular motions are consistent with the angular momentum distribution of a hypothetical uniform sphere in solid body rotation, \\begin{equation} M(<j) = M_{\\rm tot}\\left[1 - (1 - j/j_{\\rm max})^{3/2}\\right] \\label{eqt:iso} \\end{equation} where $j_{\\rm max} = \\omega R^2$ with $\\omega$ the fixed angular velocity and $R$ the radius of the sphere, and $M_{\\rm tot}$ the total mass of the sphere.  Fall \\& Efstathiou (1980) re-examined this question in the context of the more modern view according to which disk galaxies form from contracting gas within extended dark-matter halos (White \\& Rees 1978). They added the assumption that the gas and dark matter were initially well-mixed such that the distribution of $j$ in the disk is equal to that of the halo. Rather than assuming an angular momentum profile, they implicitly deduced it from the constraint that the final surface density profile of the disk is exponential. The implied disk scale radius, $r_{\\rm d}$, is related to the total pre-collapse angular momentum of the halo,  roughly $r_d \\propto \\lambda$.  Here, $\\lambda$ is the dimensionless spin parameter (e.g. Peebles 1969) \\begin{equation} \\lambda \\equiv \\frac{J\\left|E\\right|^{1/2}}{G M^{5/2}}, \\label{eq:lambda} \\end{equation} where $J$,$E$ and $M$ are the total angular momentum, energy and mass of the system, and $G$ is Newton's constant.  Angular momentum is presumably acquired by the dark matter (and gas) through tidal interactions with neighboring objects (Peebles 1969). Barnes \\& Efstathiou (1987) (following Efstathiou \\& Jones 1979; Efstathiou \\& Barnes 1983; Zel'dovich \\& Novikov 1983) utilized cosmological N-body simulations with a CDM power spectrum to make detailed predictions for halo angular momentum structure and statistics. In particular they found that the distribution of spin parameters is roughly log normal, with a median value of $\\lambda \\sim 0.05$.  They also confirmed that the angular momentum of pre-collapse halos grows roughly linearly with time, as predicted by linear tidal-torque theory (Doroshkevich 1970; White 1984). Several numerical and analtyical investigations have followed, which confirmed the nature of the $\\lambda$ distribution, addressed the alignment of the spin vectors between neighboring galaxies, the (lack of) dependence of halo spin on environment, and the effect of major mergers on halo spin  (e.g., Frenk et al. 1988; Heavens \\& Peacock 1988; Zurek, Quinn \\& Salmon 1988; Warren et  al. 1992; Catelan \\&  Theuns 1996; Cole \\& Lacey   1996; Lemson \\& Kauffmann 1997; Gardner 2000).  Armed with this global information about the initial state of the dark matter and gas mixture, several authors have used the above disk formation scenario to predict observable quantities. Blumenthal et al. (1986) addressed the rotation curves of disk galaxies in the context of the CDM hierarchical formation scenario, assuming $\\sim 10\\%$ of the mass in the gas component (Blumenthal et al. 1984).   They pointed out that the non-dissipative dark halo should react to the dissipative gas infall by considerable contraction, and showed that this process can be approximated using an adiabatic invariant.  This coupling between the halo and the disk ``conspires\" to produce a continuity between the disk-dominated and halo-dominated regions of the rotation curves, as observed (Burstein \\& Rubin 1985).  Flores et al. (1993) explored the dependence of the final rotation curves on a range of assumed parameters, including the measured spread in $\\lambda$.  Dalcanton, Spergel, \\& Summers (1997) included a more realistic initial halo density profile (Hernquist 1990) and explored the properties of both high surface brightness and low surface brightness galaxies. Mo, Mao, \\& White (1998a; 1998b, 1999) performed a similar examination within the context of various cosmologies, assuming that the dark halos followed the density profile advocated by Navarro, Frenk, \\& White (1996, 1997; NFW) and compared their predicted galaxy properties to both local and high-redshift data.  van den Bosch (1998; 2000) and  van den Bosch \\& Dalcanton (2000) extended the approach to include a bulge component and investigated the origin of the Hubble sequence of galaxy types and other aspects of galaxy formation.  The main commonality in all of these investigations is that the variation in observable galaxy properties depends almost exclusively on the initial total spin parameter of the halo hosting the galaxy. However, all of the predictions rest on the implicit assumption that, if specific angular momentum is conserved, the initial angular momentum profiles of the halos produce disks that are exponential in form.~\\footnote{An exception is Dalcanton et al. (1997), who assumed that the initial $M(<j)$ profile was that of a uniform solid body rotator, which produces disk density profiles that are similar to exponential (Ryden \\& Gunn 1987) but only over about two disk scale lengths (Olivier, Primack, \\& Blumenthal, 1991).}  Within this simple framework for disk formation, the key missing ingredient is the actual distribution of angular momentum in halos. Partially motivated by this, we set out to measure the $j$ profiles directly using a statistical sample of halos in a high resolution simulation, and to estimate the intrinsic scatter in these distributions. The long-term goal is to learn what the detailed form of halo angular momentum profiles can teach us about the origin of galaxy types, disk surface brightness profiles, and bulge properties.  An additional strong motivation for this study comes from the results of hydrodynamic simulations of disk galaxy formation in a cosmological context. Unlike the simplified, monolithic-collapse framework discussed above, the hydrodynamic treatments model gas cooling within the hierarchical growth of dark halos, and typically find that the resulting disks are significantly smaller than real galactic disks (Navarro \\& Steinmetz 2000; Weil, Eke, \\& Efstathiou 1998; Navarro \\& Steinmetz 1997; Navarro, Frenk, \\& White 1995). The problem apparently arises because most of the mass is accreted through mergers with objects whose gas component has already cooled and clumped, causing a large fraction  of their angular momentum to be transferred to the dark halo or transported outwards. Since angular-momentum transport is likely to always be from the inside out, the simple case in which $j$ is conserved during the collapse, which we can study in detail with high accuracy, can serve as a useful limiting case. This will provide a better basis for addressing the more complex physical processes required for solving the angular-momentum problem seen in the hydro simulations.  In \\se{method} we describe our methods, including the simulation and halo finding, the measurement of the spin parameter, the associated errors, and the measurement of the angular-momentum profile. In \\se{mass-dist} we discuss the measured (universal) mass distribution of angular momentum in halos. In \\se{spatial-dist} we address the spatial distribution of angular momentum, its alignment, cylindrical symmetry, and profile in spherical shells. In \\se{origin} we use both tidal-torque theory and an estimate of angular momentum  transfer due to minor mergers to explore the origin of the angular-momentum profile. In \\se{disk} we briefly describe the implications for the disk surface density profile and the ``angular momentum catastrophe\". We discuss our results and conclude in \\se{conc}. An Appendix is devoted to testing our results with a simulation having $8$ times the mass resolution, and investigating $M(<j)$ profiles of halos at high redshift. ",
        "conclusions": "\\label{sec:conc}  We have studied the angular momentum structure within dark-matter halos in the mass range $10^{12}-10^{14} \\hMsun$, using a sample of $\\sim 600$ halos in an N-body simulation of the $\\Lambda$CDM cosmology. We examined in detail the more massive of these halos, for which the angular momentum is measured more accurately, and then verified that the more typical galaxy-sized halos have a similar angular momentum distribution.  Our primary result is that the mass distribution of specific angular momentum in halos obeys a roughly universal form, which is well fit by the two-parameter function $M(<j) = \\mu \\Mvir j/(j_0 + j)$. For a fixed value of the spin parameter $\\lt$, the parameter $\\mu$ determines the shape of the distribution, with high-$\\mu$ corresponding to a pure power-law and a smaller contrast between the tails of the distribution.  The distribution of $\\mu -1$ is roughly log normal, with the $90\\%$ range spanning $\\mu - 1 = 0.06- 1.0$. Compared to an idealized uniform sphere in solid-body rotation, the simulated halos tend to have more of their mass in the tails of the distribution, especially at small $j$. The shape parameter is only weakly correlated with the spin parameter, $\\mu - 1 \\propto (\\lt)^{\\delta}$ with $\\delta = 2 \\pm 1.5$ and a correlation coefficient $r\\simeq 0.23$.  Most halos have well-aligned angular momentum vectors throughout their volume. Between 70\\% and 90\\% have alignment cosines greater than $0.7$ between the inner and outer half masses.  However, at least 10\\% of the halos have significant misalignments, which may be interesting for several reasons.  First, they tend also to be halos with low $\\mu$, indicating that the direction of {\\bf J} in accreted material, or the halo merger history, plays a role in determining the shape of the $M(<j)$ profile.  In addition, it is unlikely that the standard simple picture of disk formation could be valid within a severely misaligned halo, so perhaps these objects typically host spheroidal stellar components. Marginally misaligned halos may play a role in the formation of galactic warps (e.g., Dekel \\& Schlosman 1983).  The spatial distribution of halo angular momentum tends to be more cylindrically symmetric than spherically symmetric.  At a fixed distance from the rotation axis, mass near the equatorial plane typically has only about $\\sim 35\\%$ more specific angular momentum than corresponding mass near the poles.  The mean $j$ is spherical shells encompassing mass $M$ is well-fit by a power law, $j \\propto M^s$.  The power $s$ is distributed like a Gaussian with a mean of $s = 1.3$ and standard deviation $\\sigma = 0.3$.  We pointed out two possible explanations % for the origin of the universal angular-momentum profile. The first is based on applying a corrected version of the linear tidal-torque theory (Porciani \\& Dekel 2000) % to extended Press-Schechter mass accretion histories. The second is based on the nonlinear process of angular momentum transfer from satellite orbits in halo mergers (see also Dekel \\& Burkert 2000). % Since halo mass growth arises as a combination of merger events and relatively quiescent mass accretion, the origin of the profile in individual halos may reflect some combination of the two processes. Each of these processes seems to produce % $j$ profiles that are similar in form to those observed in simulated halos, and thus may provide a starting point for a deeper understanding of the origin of halo angular-momentum structure.  Finally, we have started to explore the implications of our universal $M(<j)$ distribution in the context of the archetypical model of galactic disk formation, namely, adiabatic baryonic infall to a rotationally supported disk while conserving angular momentum in every mass element. We find that the implied surface-density profiles, which vary as a function of both $\\mu$ and $\\lt$, deviate significantly from an exponential disk for all but the largest values of $\\mu$.  The resultant surface profiles are more extended than exponential at large $r$, and are overly concentrated at small $r$.  Since the observed light profiles of disks are closer to exponential, our result indicates that the disk formation process cannot be fully understood within this simplified picture.  We mentioned that % the general deviation from an exponential surface density at moderate radii might be overcome in this picture if a viscous transport mechanism is included (Lin \\& Pringle 1987).  As  shown by Olivier et  al. (1991), such  a mechanism tends to produce exponential stellar distributions regardless of the initial gas profile.   However, most of the expected $j$ transport typically acts to increase the central mass density, so the deviation at small radii may be a significant problem.  Even though we have assumed that the baryons experience no angular-momentum loss, we ended up with an  ``angular momentum problem'' similar to that detected in hydrodynamic simulations (e.g., Navarro \\& Steinmetz 1997, 2000, where it is associated with $j$ transport from the gas to the dark matter and from inside out).  This problem is most severe for halos with small $\\mu$, so the tendency for these halos to have misaligned angular-momentum distributions, which makes them less-likely hosts of large disk galaxies, may partly ease the problem in the limiting case where angular-momentum is conserved.  However, even if we focus on halos that tend to be well aligned, with $\\mu \\ga 1.1$, the central densities remain higher than those in exponential disks.  A solution to this problem may be to associate the derived central mass concentrations with bulges.  It is indeed possible that the central regions of disks with profiles as in \\fig{surf2} are unstable to self gravity. More detailed modeling is needed for testing the viability of this solution. In particular, one may be worried about the model over-predicting bulge to disk ratios and in particular under-predicting the number of bulge-less galaxies.  On the other hand, such an interpretation may shed new light on the very origin of the Hubble sequence, since both $\\mu$ and $\\lt$ may play a role in determining disk scale lengths, surface brightnesses, and bulge fractions."
    },
    "0011/astro-ph0011371_arXiv.txt": {
        "abstract": "Using data from the All-Sky Monitor aboard the {\\em Rossi X-ray Timing Explorer} ({\\it RXTE}), we found that the 1.5--12 keV X-ray count rate of Cyg X-1 is, on time scales from 90 seconds to at least 10 days, strongly correlated with the spectral hardness of the source in the soft state, but is weakly anti-correlated with the latter in the hard state. The correlation shows an interesting evolution during the 1996 spectral state transition. The entire episode can be roughly divided into three distinct phases: (1) a 20-d transition phase from the hard state to the soft state, during which the correlation changes from being negative to positive, (2) a 50-d soft state with a steady positive correlation, and (3) a 20-d transition back to the hard state. The pointed {\\it RXTE} observations confirmed the ASM results but revealed new behaviors of the source at energies beyond the ASM pass band. We discuss the implications of our findings. ",
        "introduction": "Cyg X-1 is identified with a binary system of $5.6$-day orbital period which contains an O$9.7$ Iab supergiant and a compact object that is believed to be a black hole (\\cite{bolton72}; \\cite{webster72}). Observations  indicate that the system usually assumes one of the two states, the hard state and the soft state (\\cite{oda77}; \\cite{nolan84}; \\cite{tanaka95}).  Most of the time, Cyg X-1 stays in the hard state, where its soft X-ray (often $2$--$10$ keV) flux is relatively low and the X-ray spectrum is hard.  Every few years, the system undergoes a transition from the hard state to the soft state,  during which the soft X-ray flux increases, often by a factor of more than $4$, and the X-ray spectrum softens. It remains in the soft state for weeks to months before returning to the hard state. There is a strong anti-correlation between the soft and the hard (e.g., $>$ 20 keV) X-ray  flux. Consequently, the bolometric X-ray luminosity does not vary significantly (\\cite{zhang97a}). The transition between the two states lasts from less than a day to more than a week.  A number of models have been proposed to explain the spectral evolution of Cyg X-1 during the state transitions. For instance, Ichimaru (1977) suggested that the physical condition of the accreted gas near the disk outer-boundary could drive the disk into either an optically thick state or an optically thin state, which correspond to the soft and hard state respectively. Zhang, Cui, \\& Chen (1997) argued, based on the effects of black hole rotation, that the state transition of Cyg X-1 may be caused by a temporary reversal of the disk from being retrograde (hard state) to prograde (soft state), which can occur in wind accretion systems (e.g., \\cite{matsuda87}; \\cite{ruffert97}). In the magnetic flare model (\\cite{matteo99}), the soft (hard) state corresponds to a lower (higher) scale height of magnetic flares above the accretion disk. The flares are energized by the reconnection of magnetic flux tubes rising from the accretion disk due to magnetic buoyancy instability.  In the soft state, intense flares close to the disk greatly enhance the soft photon field which results in a soft X-ray spectrum. In the hard state, the flare is triggered high above the disk, the system is ``photon starved'' and thus results in a hard, Comptonized spectrum. In the framework of advection-dominated accretion flows (ADAFs), Esin et al. (1998) argued that the spectral states of Cyg X-1 are uniquely determined by the mass accretion rate \\.{\\it m}. In the hard state, \\.{\\it m} is relatively low, the inner edge of the thin disk is far away from the black hole and thus the emission from the disk is weak compared to that from the large, optically thin ADAF region. In the soft state, the \\.{\\it m} is higher, which causes the ADAF region to shrink and the thin disk to extend closer to the black hole. Therefore, at low energies the emission from the disk dominates over that from the smaller ADAF region. More data are needed to distinguish these models.  In this Letter, we report the results from a quantitative study of the correlation between the X-ray flux and the spectral properties of Cyg X-1, on  time scales of 90 s to at least 10 days, in both the hard and soft states, as well as during the transition phases. ",
        "conclusions": "The ASM results show that the transition phases lasted for about 20 days during the 1996 state transition of Cyg X-1, as opposed to $\\lesssim$ 7 days as indicated by the change in the soft X-ray flux (see Fig.~1).   That is, for about 20 days at the beginning and near the end of the state transition episode, the system was  in a transitional process even though the soft flux was generally high. A similar conclusion was drawn by Cui et. al (1997a, 1997b) based on the evolution of the power density spectra (PDS).   We therefore suggest that the spectral states of Cyg X-1 are better defined by the correlation between the soft X-ray flux and the spectral hardness of the source than by the soft flux alone.  The PCA results confirm our findings based on the ASM data. Furthermore, they reveal that the orbital effects could account for the observed negative correlations on time scales less than 5.6 days, for time periods sufficiently far away from the soft flare or the state transition. This is consistent with the fact that there is a broad absorption-like dip in the orbital light curves obtained from the ASM data, likely caused by absorption and scattering of the X-rays by the stellar wind (e.g., Wen et al. 1999).  Our best-fit model ASM orbital light curve and hardness ratios (without noise) would yield $r_{s1}$ to be around -0.6 with a 0.1 day time bin in a 5.6 day correlation interval. The observed correlation strength is weaker ( $r_{s1} \\sim -0.2$), probably reduced by noise.  On the other hand, the observed positive correlation in the soft state is the opposite to what we would expect from the orbital effect.  This is consistent with our previous conclusion (Wen et al. 1999) that orbital modulation is much smaller (if present at all) in the soft state.    Both the ASM and the PCA data indicate the existence of the flux-hardness anti-correlation in the hard state on time scales longer than the $5.6$ day orbital period.  Moreover, the hard state data near the state transition episode and during the soft flare both contribute to this observed anti-correlation in a similar fashion (Fig.~\\ref{pca}).  The PCA spectral fitting indicates that these observations also share similar spectral properties, which deviate somewhat from that of typical hard state data (see data section). This seems to lend support to the notion that soft flares are ``failed'' state transitions. Similar soft flares occur randomly on time scales of months to years in the ASM light curve. We therefore conclude that the anti-correlation in the hard state on time scales longer than the orbital period is intrinsic to the source and is probably related to the mechanism that causes the state transition.  Models employing a simple disk--corona geometry predict that, as the mass accretion rate through the disk increases, the soft X-ray flux increases, which provides more seed photons for the inverse-Comptonization process and thus cools the hot electrons in the corona. This would be manifested observationally as the ``pivoting'' of the energy spectrum of the source. Such spectral pivoting is known to occur in Cyg X-1 during a state transition (e.g., \\cite{nolan84}; \\cite{zhang97a}), and the pivoting energy is in the range of 10--20 keV. This may explain the observed flux-hardness anti-correlation at low energies (below the pivoting energy) for the state transition (and perhaps soft flares). The same might also be true for the general long-term evolution of the source in the hard state, if the pivoting phenomenon is universal (but more pronounced during a state transition). Interestingly, this may also explain the apparent lack of correlations around 12--25 keV (Fig. \\ref{pca}), since the pivoting energy is right in this energy range. Clearly, the same scenario cannot be applied to the soft state, where the correlation is observed to be (strongly) positive, unless the pivoting energy has moved to a very low energy (below the ASM pass band). Such a reduction in the pivoting energy is not apparent in our PCA data and is inconsistent with the lack of correlations in the 12--25 keV band. It is, therefore, likely that different physical processes are involved in the soft state (see \\cite{zdziarski00} for a follow-up investigation of possible models).     Li, Feng, \\& Chen (1999) discovered similar correlations between the count rate and hardness for Cyg X-1, using the PCA data, but on much shorter time scales (0.01--100 s) and with little coverage in the hard state. In their study, they chose two energy bands: 2--6 keV and 13--60 keV, from which the hardness ratio was derived. Such a choice of energy bands unfortunately masks the difference in the correlations between the count rate and the spectral shape above and below 12 keV for the soft state. Using the BATSE data (20-200 keV), on the other hand, Crary et al. (1996) did see a  lack of correlation between the 45--140 keV energy flux and the photon index within both data groups of high and low flux, which were presumed to correspond to the hard and soft states, respectively. However, they did not have the necessary soft-X-ray data to see the count rate-hardness correlation that we found.  In this regards,  our work bridges critical gaps in  those two investigations and provides new insight into the overall spectral behavior of this system."
    },
    "0011/astro-ph0011147_arXiv.txt": {
        "abstract": "Following the recent measurement of the acoustic peak by the BOOMERanG and MAXIMA experiments in the CMB anisotropy angular power spectrum, many analyses have found that the geometry of the Universe is very close to flat, but slightly closed models are favoured. In this paper we will briefly review how the CMB anisotropies depend on the curvature, explaining any assumptions we could make and showing that this skewness towards closed models can be easily explained by degeneracies in the cosmological parameters. While it is difficult to give independent constraints on the cosmological constant and/or different forms of dark energies, we will also show that combining CMB measurements with other observational data will introduce new and tighter constraints, like $\\Omega_\\Lambda > 0$ at high significance. ",
        "introduction": "In the most general and simple inflationary scenario \\cite{guth} the overall present energy density of the universe must be equal to the so-called critical energy density ($\\Omega=\\rho / \\rho_c =1$, with $\\rho_c=3H_0^2/8 \\pi G$). In fact, the condition for inflation ($\\ddot a_{inf} >0$, and so $d / dt (H_{inf}^{-1}/a_{inf}) < 0$) is precisely that which drives $\\Omega(t)$ towards $1$ in the Friedmann equation  \\bee \\Omega(t)-1={k \\over {a^2(t)H^2(t)}} \\ee  during the inflationary period. This prediction, however, taken with the {\\it standard} CDM model of structure formation, is apparently in disagreement with a combined set of observations, such as density-velocity galaxy field comparisons (\\cite{davis}), pairwise galaxy velocities (\\cite{jusk}), X-ray clusters temperature function evolution (\\cite{eke}, \\cite{henry}) and velocity dispersions \\footnote{A more conservative approach would say that the situation is rather unclear and the X-ray clusters data and theoretical modelling can be sufficiently large to prevent an unambigous exclusion of $\\Omega_m=1$ (see {\\it e.g.} \\cite{cola}, \\cite{viana}, \\cite{borg}, \\cite{blan})}(\\cite{carl})], which point in favor of a low density universe \\footnote{From here on, with $\\Omega$ we will indicate only the {\\it present} value} ($\\Omega_{matter} <1$). In order to solve the discrepancy, open ($\\Omega=\\Omega_{matter} <1$) inflationary models have been proposed (\\cite{sasa}, \\cite{buche} \\cite{linde}). This type of model highlights limitations on the predictiveness of the inflationary scenario, which is supposed to have the advantage of removing any dependence on initial conditions from our present observable universe. However, even if this picture leads to a more complicated phenomenology, it generally determines $\\Omega$ directly from parameters of the physical theory.  On the other hand, another way to keep low-density models compatible with the simplest model of inflation is to introduce a cosmological constant, $\\Lambda$, such that $\\Omega\\equiv\\Omega_{matter}+\\Omega_{\\Lambda}=1$. The presence of such a cosmological constant, which is compatible with all the above observations since they are practically insensitive to it, is also preferred by measurements of the magnitude-redshift relationship in high-redshift type Ia supernovae (\\cite{perlm}, \\cite{riess}). Nonetheless, this {\\it 'natural'} solution introduces the cosmological constant problem (\\cite{weinb1}) that is perhaps even more acute in inflationary cosmology (see \\cite{bran} for a review). Thus, an accurate determination of the {\\it present} overall density parameter $\\Omega$, even if not a {\\it panacea} for the cosmological scenario, is at least extremely important in understanding which theoretical framework could explain the above conflicts.  There is much experimental evidence for the presence of a peak in the CMB angular power spectrum (\\cite{mille}, \\cite{mausk}).  Furthemore, with the recent release of the BOOMERanG-98 \\cite{debe} and MAXIMA-1 \\cite{hano} spectra, the shape and position of the peak has been detected with unprecedented accuracy. This result, apart from being in wonderful agreement with the standard scenario of primordial {\\it adiabatic} fluctuations, has important consequences on the parameter $\\Omega$ (\\cite{melk00}, \\cite{lange}, \\cite{balbi}, \\cite{jaffe}, \\cite{ruth}). As previously noted and already well explained in the literature, the effect of the curvature is to change the relationship between the physical scales on the Last Scattering Surface (LSS) and the corresponding angular scales (see., {\\it e.g.} \\cite{weinb}, \\cite{turn}). In an open universe, for example, the geodetics focalize in such a way that a particular angular scale will correspond to a greater physical scale on the LSS than the one expected in a flat model. The immediate result is a shift in the radiation angular power spectrum (the so-called $C_\\ell$'s), and thus a dependence of the position of the first peak $\\ell_{peak}$ on the curvature (\\cite{wils}, \\cite{kamion}, \\cite{silk}, \\cite{hu}, \\cite{scott}). The CMB power spectrum is therefore a powerful tool for the determination of the curvature and so of the overall energy density.  \\begin{figure}[htb] \\includegraphics[width= 10cm,height=8cm,angle=-90]{melchiorria_fig0.ps} \\caption{The BOOMERanG and MAXIMA $C_{\\ell}$'s data together with the most viable open, closed and flat models from present non--CMB observations.} \\label{figregion} \\end{figure} ",
        "conclusions": "The BOOMERanG and MAXIMA data support the main prediction of the inflationary paradigm: that the geometry of the universe is flat. The small deviations towards closed models reported in various analyses (\\cite{lange},\\cite{jaffe}) can be easily explained by the degeneracies in the cosmological parameters, which make {\\bf more closed} models compatible with the data. These conclusions are considerably strengthened by the inclusion of other cosmological data such as measurements of the Hubble constant, the overall matter density $\\Omega_m$ and the accelerating expansion rate indicated by observations of distant Supernovae. At the same time, $\\Omega_\\Lambda$ and other forms of 'dark energy' cannot be well determined by the CMB data alone, in spite of their high precision. This does not mean that CMB measurement are not useful in the determination of such parameters: combining the CMB data with constraints from observations of large-scale-structure and from observations of SN-Ia increases the extent to which $\\Lambda$ can be quantified, with $\\Omega_{\\Lambda}>0$ at $\\sim 7 \\sigma$. Future data from the PLANCK and/or SNAP satellites will hopefully enable us to resolve the 'dark energy' puzzle."
    },
    "0011/astro-ph0011417_arXiv.txt": {
        "abstract": "The feasibility of using near-infrared observations to discover supernovae in the nuclear and circumnuclear regions of nearby starburst galaxies is investigated.  We provide updated estimates of the intrinsic core-collapse supernova rates in these regions.  We discuss the problem of extinction, and present new estimates of the extinction towards 33 supernova remnants in the starburst galaxy M~82.  This is done using H~I and H$_{2}$ column density measurements.  We estimate the molecular to atomic hydrogen mass ratio to be 7.4 $\\pm$ 1.0 in M~82.  We have assembled near-infrared photometric data for a total of 13 core-collapse supernovae, some unpublished hitherto.  This constitutes the largest database of IR light curves for such events.  We show that the IR light curves fall into two classes, ``ordinary'' and ``slow-declining''.  Template $JHKL$ light curves are derived for both classes.  For ordinary core-collapse supernovae, the average peak $JHKL$ absolute magnitudes are --18.4, --18.6, --18.6, and --19.0 respectively.  The slow-declining core-collapse SNe are found to be significantly more luminous than the ordinary events, even at early times, having average peak $JHKL$ absolute magnitudes of --19.9, --20.0, --20.0, and --20.4 respectively.  We investigate the efficiency of a computerised image subtraction method in supernova detection.  We then carry out a Monte Carlo simulation of a supernova search using $K$-band images of NGC~5962.  The effects of extinction and observing strategy are discussed.  We conclude that a modest observational programme will be able to discover a number of nuclear supernovae. ",
        "introduction": "Core-collapse supernovae (SNe) are observed to occur in sites of recent star formation.  Such regions contain large quantities of dust, especially the nuclear ($r$ $\\la$ 500~pc) regions of starburst galaxies.  Consequently, SN search programmes working at optical wavelengths probably miss a significant number of events in starburst galaxies due to dust obscuration.  Thus, the observed supernova rates may be only a lower limit for the true supernova rate. The inferred star formation rates and radio observations of young supernova remnants (SNRs) indicate that highly-extinguished supernovae should exist in the nuclear regions of nearby starburst galaxies.  In NGC~253 ({\\it e.g.} Ulvestad $\\&$ Antonucci 1997), NGC~2146 (Tarchi {\\it et al.}  2000), M~82 ({\\it e.g.} Allen $\\&$ Kronberg 1998) and NGC4038/39 (Neff \\& Ulvestad 2000), several compact non-thermal radio sources have been identified as young SNRs indicating explosion rates of obscured supernovae of around 0.1~yr$^{-1}$.  Also, high resolution radio observations have recently revealed a group of luminous radio supernovae in the nuclear regions of ARP~220 (Smith {\\it et al.} 1998). Yet only one supernova event (SN 1940E in NGC 253) has ever been directly observed in (or in front of) the nuclear regions in any of these five galaxies.  It is desirable to discover and study SNe in the centres of nearby starburst galaxies.  In addition to providing a better estimate of the core-collapse rate in the local universe, such a survey would offer an opportunity to study the behaviour of supernovae in the dusty and dense nuclear environment and to probe the nuclear extinction. Indeed, SNe exploding in the nuclear regions of galaxies may differ significantly from ordinary SNe due to the higher density environment as suggested in the starburst model for AGN (Terlevich {\\it et al.}  1992).  The best-studied example of a core-collapse SN exploding within a high density circumstellar environment is SN 1998S (Fassia {\\it et al.} 2000a, Fassia {\\it et al.} 2000b, Gerardy {\\it et al.} 2000, Leonard {\\it et al.} 2000).  Ultimately, such a survey would also be valuable in guiding attempts to determine the rates of high redshift core-collapse supernovae.  Such a high redshift study is one of the aims of the NGST (New Generation Space Telescope) (Jorgensen {\\it et al.} 1997; Madau {\\it et al.}  1998, Dahlen \\& Fransson 1999; Sullivan {\\it et al.} 2000).  However, before starting to investigate the evolution of the SN rate with redshift it is important to have an accurate measurement of the local SN rate and to understand the behaviour of supernovae within the dusty, high density starburst regions of galaxies.  The feasibility of searching for obscured supernovae in nuclear starburst regions has been discussed in two papers. Van Buren \\& Norman (1989) considered the use of mid-IR narrow-band (10.52 $\\mu$m) imaging.  They estimated that supernovae would be observable in the mid-IR with ground-based 4-10 m telescopes at distances of 15-40 Mpc in one night's integration. Supernova detection in the $K$-band was studied by Grossan {\\it et al.} (1999). They emphasised the need for high spatial resolution to allow successful image subtraction in the nuclear regions of galaxies.  An optical search for supernovae in a sample of 142 nearby star-forming galaxies was carried out in 1988-1991 by Richmond {\\it et al.} (1998).  The observations were carried out with a $500\\times500$ pixel CCD camera (0.58\"/pixel) on a 1~m telescope at the Lick Observatory.  Two search procedures were used.  In the first method, the search images were compared by eye.  This yielded five supernovae, all of which were outside the host galaxies' extinguished nuclear regions.  The deduced SN rates were similar to those measured in normal galaxies.  In the second method, differential photometry was carried out on the nuclei of the galaxies.  However, they did not discover any brightening attributable to nuclear supernovae.  From this, they deduced upper limits for the {\\it unobscured} SN rates within the nuclear regions {\\it viz.} $<$ 9h$^{2}$~SRU\\footnote{ One SRU is the number of supernovae per century per 10$^{10}$L$_{\\odot}$ galaxy blue luminosity, and H$_{\\rm 0}$ = h $\\times$ 100 kms$^{-1}$/Mpc.}, $<$ 12h$^{2}$~SRU, and $<$ 7h$^{2}$~SRU for Types Ia, Ib/c, and II, respectively.  This work confirmed that, given the likely high obscuration in nuclear starbursts, for a search for supernovae in such regions to be successful it would have to be carried out at infrared wavelengths.  Van Buren {\\it et al.} (1994) conducted a $K$-band survey for supernovae in starburst galaxies at NASA's Infrared Telescope Facility, Hawaii for about 2 years. Most of the observations were made with PROTOCam, and the Richardson-Lucy image restoration algorithm was used to enhance the resolution of the survey images.  They discovered SN~1992bu in NGC~3690 in $K$-band images$\\footnote{ www.ipac.caltech.edu/ipac/info/sn1992/sn1992\\_NGC3690.html}$.  The supernova is located $\\sim$6\" from Core B1 of the galaxy. They measured $K$-band magnitudes of +16.6, +17.2, and +18.1 in three images separated by about a month.  No other observation of this supernova was reported, and thus it is not possible to say if it was located within the obscured nuclear regions.  A more recent (1992-1994) near-infrared (NIR) search for supernovae in 177 nearby IRAS galaxies has been carried out in the $K'$-band by Grossan {\\it et al.} (1999). The observations were performed with the 2.3~m Wyoming IR Observatory (WIRO) telescope using the Michigan IR camera (MIRC), with a 2.2\"/pixel resolution and 128$\\times$128 pixels. The seeing was typically between 0.7'' and 1.0\" and so the seeing disk was usually undersampled. The length of the time between two observations of a given galaxy varied typically from 1 to 3 months. They did not discover any supernovae during the period, limiting the SN rate {\\it outside the nuclear regions} ($>$ 15\" radius) to less than 3.8 FIRSRU$\\footnote{ One FIRSRU is the number of supernovae per century per 10$^{10}$L$_{\\odot}$ galaxy far-IR luminosity (see Section~2).}$ in their sample galaxies within 25~Mpc distance.  They concluded that this negative result was due to the poor resolution of the camera and that a higher resolution NIR supernova search covering only the inner $\\sim$450pc of each galaxy would be more productive.  The latest reported attempt to detect SN explosions in the nuclear regions of starburst galaxies was by Bregman {\\it et al.} (2000). They observed a sample of 10 galaxies within 7~Mpc distance using ISOCAM with a pixel size of 3''. After continuum subtraction they looked for traces of [Ni~II] 6.63 $\\mu$m line emission which would have indicated recent SN explosions.  However, they did not detect any [Ni~II] emission in the sample galaxies.  They presented an upper limit for the SN rate in their galaxy sample normalised to M~82 viz. 0.3~yr$^{-1}$. Given the far-IR luminosity of M~82 (Appendix~A) this corresponds to 7.2 FIRSRU.  Although, as argued above, the discovery of such SN events would clearly be important, given the uncertainties in {\\it e.g.} nuclear extinction or SN magnitudes it is less obvious what might be achieved in practice with existing telescopes.  We have therefore carried out a study to test the feasibility of discovering obscured supernovae in the nuclear regions of starburst galaxies.  In this paper we report the results of this study.  In Section~2, we review indirect methods of assessing the nuclear supernova rate and provide an updated estimate of this parameter.  In Section~3, extinction towards nuclear supernovae is discussed.  In particular, we present a new estimate of the extinction and its distribution in the starburst galaxy M~82.  In Section~4, we present $JHKL$-band light curves for a total of 13 events, some unpublished hitherto.  This constitutes the most complete database of core-collapse SN IR light curves ever published.  We show that the IR light curves fall into two classes, ``ordinary'' and ``slow-declining''.  Template $JHKL$ light curves are derived for both classes.  In section 5, supernova detection using an image subtraction method is described and tested, and in section~6 Monte Carlo simulations are used to test the feasibility of a practical supernova search.  The results are summarised in Section~7. ",
        "conclusions": "The feasibility of detecting supernovae in the nuclear regions of nearby starburst galaxies has been investigated.  This has entailed the examination of expected SN rates, extinction, and SN magnitudes in such regions.  An intrinsic SN rate of 2.7 $\\times$ 10$^{-12}$ $L_{\\rm FIR}$ / $L_{\\odot}$~~yr$^{-1}$ was estimated.  Extinction values ranging between $A_{K}$ = 2 and 4 were judged to be typical. A particularly detailed examination of the M~82 extinction distribution was carried out. A mean $A_{V}$ of 24 was obtained from hydrogen column density measurements.  Template core-collapse SN light curves in $JHKL$ were assembled. Two types were identified - ordinary and slow-declining.  The slow-decliners were found to be significantly more luminous than the ordinary core-collapse SNe, even at early times.  Using the Optimal Image Subtraction method, it was found that supernovae outside the innermost $\\sim$1'' of a galaxy can be easily detected in $K$-band images.  Imaging of a sample of 45 galaxies on a 6 month cycle, would probably reveal at least 1 SN per cycle.  However, the actual number of detections could be significantly higher if slow-decliners are more common in nuclear regions.  Near-infrared observations of SNe in nuclear starbursts will provide us with a valuable new probe of the conditions and star formation rates in the obscured starburst regions of galaxies."
    },
    "0011/astro-ph0011284_arXiv.txt": {
        "abstract": "Gamma-ray burst sources with a high luminosity can produce electron-positron pair cascades in their environment as a result of back-scattering of a seed fraction of their original hard spectrum. The resulting spectral modifications  offer the possibility of diagnosing not only the compactness parameter of the $\\gamma$-ray emitting region but also the baryonic density of the environment external to the burst. ",
        "introduction": "The non-thermal $\\gamma$-ray spectrum of Gamma-Ray Burst sources (GRBs) is thought to arise in shocks which develop beyond the radius at which the relativistic fireball from the initial event has become optically thin to scattering.  However, the observed spectra are hard, with a significant fraction of the energy above the $\\gamma\\gamma \\to \\epm$ formation energy threshold, and a high compactness parameter can result in new pairs being formed outside the originally optically thin shocks responsible for the primary radiation. New pairs can be made as some of the initial energetic photons are backscattered and interact with other incoming photons. Such effects have been considered by Madau \\& Thompson (2000) and Thompson \\& Madau (2000), who investigated the acceleration of new pairs for a particular fireball model. Dermer \\& B\\\"ottcher (2000) considered the effect of pair formation for an external shock model of GRB,  in which the compactness parameter is relatively low, while Madau, Blandford \\& Rees (2000)  investigated the effects of Compton echos produced by pairs.  Here we present a simplified discussion of the generic effects of pair formation in a specific class of models. We suppose that the central engine gives rise to an unsteady baryonic wind, which is relativistic, carries a magnetic field, and lasts for a time $t_w$.  A fraction of the wind energy is converted into  $\\gamma$-rays  via internal shocks; the remaining wind energy drives a blast wave which decelerates as it sweeps up the external medium, and gives rise to the afterglow emission. The  $\\gamma$-rays would propagate ahead of the blast wave, leading to pair production (and an associated deposition of momentum)  into the external medium.  The pair cascades saturate after the external (pair-enriched) medium reaches a critical bulk Lorentz factor, which is generally below that of the original relativistic wind. For external baryonic densities similar to those in molecular clouds the pairs can achieve scattering optical depths $\\tau_\\pm \\siml 1$. Even for less extreme external densities the effect of the additional pairs can be substantial, increasing the radiative efficiency of the blast wave  and leading to distortions of the original spectrum. This provides a potential tool  for diagnosing the compactness parameter of the bursts and thus the radial distance at which shocks producing an observed luminosity can occur. It also provides a tool for diagnosing the baryonic density of the external environment of the bursts, and testing the association with star-forming regions. ",
        "conclusions": "We have discussed, in the context of a standard internal/external shock model of gamma-ray bursts (which are normally assumed to occur after the original fireball has become optically thin) the  $\\gamma\\gamma\\to \\epm$ cascades triggered by the back-scattering of seed gamma-ray photons on the external medium. This effect can modify the initial scattering optical depth of the outflow at radii $r\\siml r_\\ell \\simeq 4\\times 10^{15}\\L50^{1/2}\\tw1^{1/2}~\\hbox{cm}$ (equation [\\ref{eq:rell}]), which is comparable to the radii of external shocks of equation (\\ref{eq:rdec}) at which the afterglow begins, and is generally larger than the typical internal shock radii given by equation (\\ref{eq:rint}).  The spectral effects of the pairs  on the burst and the afterglow can be substantial, and within radii $\\sim r_l$ they can affect the dynamics.  Pair production can increase  the optical depth outside of the shocks by  up to  $\\siml 10^5$;  self-shielding ensures that the maximum scattering optical depths achievable by the pairs is $\\tau_\\pm\\sim 1$. For typical interstellar densities the pair opacity $\\tau_\\pm \\ll 1$, which does not significantly affect the gamma-ray spectrum. The number of pairs may nonetheless be large enough to increase the radiative efficiency and soften significantly the radiation spectrum of the afterglow reverse shock, where the same energy is shared among a number of $\\epm$ which can exceed that of the original $e^-$ and $p^+$ of the ejecta.  The  pair production processes themselves (determining $r_c,~ r_l$, equations [\\ref{eq:rcpp}], [\\ref{eq:rell}]) just depend on the ``seed\" $\\gamma$-ray photon flux (which are here postulated to come from internal shocks).  The manifestations depend on the external density and on the initial dimensionless entropy or bulk Lorentz factor $\\eta$. The external baryon density $n_{ext}$ determines the optical depth that can be built up through back-scattering and pair multiplication. This affects whether the pair optical depth gets up to unity, with smearing and reprocessing of the primordial $\\gamma$-ray spectrum, or whether it merely makes the blast wave more radiative. Madau and Thompson (2000) have made this point, in the context of a specific fireball model.  The external density (along with the initial Lorentz factor $\\eta$) determines when the  outer shock and the reverse shock become important and whether this happens within the radius already polluted with pairs  (and pre-accelerated by radiation pressure before the shock hits). There are two rather different cases depending on whether or not $\\eta^2$ is less than $r_l/ct_w$). In the former case the external shock responsible for the afterglow occurs beyond the region ``polluted\" by new pairs, while in the second case the afterglow shock may experience, after starting out in the canonical manner, a ``resurgence\" or second kick as its radiative efficiency is boosted by running into an $\\epm$-enriched gas.  Internal shocks in the pair-dominated external plasma  can lead to self-absorbed radiation at $\\sim 10^9-10^{10}$ Hz, while the swept-up pairs can also contribute  a $10^{11}-10^{12}$ Hz `prompt' signal, which precedes the onset of the standard deceleration afterglow phase.  Additional effects are expected when $\\tau_\\pm \\to 1$. This requires external baryon densities at radii $r<r_\\ell$ of $n_p \\simg n_{c,p} \\sim 10^5 \\L50^{-2/5} \\tw1^{-3/5} \\cmcui$. Such high densities would only be expected if the burst is associated with a massive star in which prior mass loss  led to  a dense circumstellar envelope  The pair optical depth saturates to $\\tau_\\pm\\sim 1$ and in addition to an increased efficiency and softer spectrum of the afterglow reverse shock, the original gamma-ray spectrum of the GRB will be modified as well.  The specific nature of this spectral modification depends on the value of the luminosity, which influences (equation [\\ref{eq:rcpp}]) the bulk Lorentz factor of the reprocessing pair cloud before it has been swept up by the ejecta, and also on the extent to which the outflow is beamed. One of the consequences of such a critical external density leading to $\\tau_\\pm\\sim 1$ would be the presence of an X-ray quasi-thermal pulse, whose total energy may be a few percent of the total burst energy. In the case of even more extreme densities, there are other interesting possibilities. For instance, dense blobs of Fe-enriched thermal plasma would emit strong recombination features, as well as annihilation radiation features, if the normal electron density were augmented by extra pairs.  Even in isotropic situations, the spectrum would be modified  by transmission through a pair plasma of optical depth unity. The effect is maximal for photons of energy $m_e c^2$ in the frame of the pair plasma: for higher energies, the Klein-Nishina cross-section is smaller; for lower energies, the scatterings are almost elastic.  If, for instance, $\\Gamma_c\\sim 40$, the gamma-ray spectrum around photon energies $\\eps\\sim \\Gamma_c m_e c^2 \\sim 20$ MeV  would be depressed by a factor $\\sim 1$ relative to its original value, smoothly rejoining its original value above and below that energy. For a beamed primary output, however, there would be a suppression at lower energies (where the scattering is in the Thompson regime) because the scattered photons would be spread over a wider angle. Pair-induced processes would therefore yield evidence on the beaming properties of the bursts.  Irrespective of the external density, the processes discussed here suggest that bursts and afterglows may have a more complex spectrum and time-structure than 'standard' models suggest. But the effects are especially interesting when the external density is high: they probe the environment of GRBs, and thus can offer clues to the nature of the progenitor  stars,  and their location within the host galaxy. For instance, a quasi-thermal X-ray pulse accompanying the gamma-ray emission could be indicative of an external circum-burst density of at least $10^4-10^5~\\cmcui$. While quasi-thermal X-ray pulses might also arise due to other reasons, e.g. from an underlying optically thick central engine, if the X-ray luminosity scales as equation (\\ref{eq:xrprec}) and is accompanied by radio or far-IR signals such as in equations (\\ref{eq:nusyr})(\\ref{eq:nusyf}), this could be indicative of birth in a dense environment from a massive progenitor.  Detailed Monte Carlo simulations (Ramirez-Ruiz \\etal, 2001) should provide a more detailed assessment of the self-consistent spectrum of a GRB in the presence of self-induced pair formation."
    },
    "0011/astro-ph0011551_arXiv.txt": {
        "abstract": "We consider envelopes of DB white dwarfs which are not composed of pure $^4$He, but rather a mixture of $^{3}$He and $^4$He. Given this assumption, the same diffusive processes which produce a relatively pure H layer overlying a He layer in the DA's should work to produce a $^{3}$He layer overlying a $^4$He layer in the DB's.  We examine the relevant timescales for diffusion in these objects, and compare them to the relevant evolutionary timescales in the context of the DBV white dwarfs.  We then explore the consequences which \\he{3} separation has on the pulsational spectra of DBV models. Since GD~358 is the best-studied member of this class of variables, we examine fits to its observed pulsation spectrum. We find that the inclusion of a \\he{3} layer results in a modest improvement in a direct fit to the periods, while a fit to the period {\\em spacings} is significantly improved. ",
        "introduction": "The two greatest successes of white dwarf asteroseismology have come from the analysis and interpretation of Whole Earth Telescope (WET) data on the objects GD~358 and PG~1159 \\citep{Winget91,Winget94}. The temporal spectra of both objects show well-defined multiplet structures as well as many consecutive radial orders of the same $\\ell$. This allows us immediately to estimate the rotational frequency, $\\Omega$, and the mean period spacing, $\\langle \\Delta P \\rangle_{\\ell}$, with the latter quantity putting strong constraints on the mass of the best-fitting model. Furthermore, the variations from equidistant period spacing, i.e., ``mode trapping'', give us information (at least in the case of GD~358) about the radial structure of the star, such as the thickness of its surface He layer.  In a detailed analysis of the frequency distribution in GD~358, \\citet{Bradley94a} found evidence of a chemical transition zone located at a mass depth of $1.5 \\times 10^{-6} \\Mstar$, which they interpreted as the C/He boundary. A He layer thickness of $\\sim$10$^{-6} \\Mstar$ is thinner than expected on naive evolutionary grounds, as well as thinner than suggested for the DA white dwarfs by the asteroseismological work of \\citet{Clemens93}.  While the fit of Bradley \\& Winget explained the trapping features in the neighborhood of 700 sec reasonably well, it did not do as well for periods near 500 sec. They found evidence that these periods near 500 sec might be better fit if an additional chemical transition zone at a depth of $\\sim$10$^{-2} M_{\\star}$ is assumed.  Independent of the above considerations, if the theory of diffusion of chemical species in stars is correct, then at least a partial separation of different isotopes of the same element must occur in white dwarfs, as was pointed out to one of us by Clayton (1988, private communication). At the time, there were no known observational consequences; this has since changed with the advent of white dwarf asteroseismology.  While the Galactic number ratio of \\he{3} to \\he{4} has been measured to be of order $10^{-4}$ in many astrophysical environments \\citep{Galli95,Prantzos96}, this does not necessarily hold for the stellar cores which are presumably the white dwarf progenitors. In fact, standard evolutionary theory would suggest quite low levels of \\he{3} in these objects \\citep{Galli95}. Our view is that given the theoretical uncertainties in the late stages of stellar evolution, we should not rule out any possibilities, and indeed should seek to make as many independent measurements as possible.  These considerations led to two independent motivations for considering \\he{3} diffusion \\citep{Winget98}.  First, it is a process which may be a generic feature of white dwarf cooling, and, as such, should be considered as part of our asteroseismological analyses.  Second, it might allow us to place the mode trapping data for GD~358 in a different context: this star could have a C/\\he{4} transition zone at $\\sim 10^{-2} \\Mstar$ and a thinner \\he{4}/\\he{3} transition zone at $\\sim 10^{-6} \\Mstar$. If this is the case, then the DBV's and the DAV's would again have the same order of magnitude He layer thicknesses. To determine the plausibility of these hypotheses, we first examine the relevant diffusion timescales, and then the effect which such a layering structure has on our fits of GD~358. ",
        "conclusions": "From this preliminary analysis, we have shown that diffusion theory applied to white dwarfs predicts that any \\he{3} initially present in the DBV's down to a depth of $\\sim$10$^{-4} \\Mstar$ should have diffused upward to produce a surface layer of \\he{3}. In addition to being detectable spectroscopically (Koester, private communication), such a \\he{3} layer can significantly affect the asteroseismological fits,  and therefore needs to be included in such analyses.  We find that although the inclusion of a \\he{3} layer results in only a marginal improvement to the fits to the periods, the fit to the period {\\em spacings} is significantly improved. This is because the period spacings are more sensitive to the fine structure which a \\he{3}/\\he{4} transition zone produces than the periods themselves are.  Finally, we find that an O/C core (essentially, the transition in the core from an O/C mixture to pure C) fits the observed pulsational spectrum of GD~358 {\\em much} better than a pure C core, in agreement with \\citet{Metcalfe00}, who first obtained this result. This gives us the hope of someday being able to constrain the prior nuclear burning history of GD~358 and other pulsating white dwarfs."
    },
    "0011/gr-qc0011028_arXiv.txt": {
        "abstract": "A quasi-equilibrium (QE) computational scheme was recently developed in general relativity to calculate the complete gravitational wavetrain emitted during the inspiral phase of compact binaries.  The QE method exploits the fact that the the gravitational radiation inspiral timescale is much longer than the orbital period everywhere outside the ISCO. Here we demonstrate the validity and advantages of the QE scheme by solving a model problem in relativistic scalar gravitation theory.  By adopting scalar gravitation, we are able to numerically track without approximation the damping of a simple, quasi-periodic radiating system (an oscillating spherical matter shell) to final equilibrium, and then use the exact numerical results to calibrate the QE approximation method. In particular, we calculate the emitted gravitational wavetrain three different ways: by integrating the exact coupled dynamical field and matter equations, by using the scalar-wave monopole approximation formula (corresponding to the quadrupole formula in general relativity), and by adopting the QE scheme.  We find that the monopole formula works well for weak field cases, but fails when the fields become even moderately strong. By contrast, the QE scheme remains quite reliable for moderately strong fields, and begins to breakdown only for ultra-strong fields.  The QE scheme thus provides a promising technique to construct the complete wavetrain from binary inspiral outside the ISCO, where the gravitational fields are strong, but where the computational resources required to follow the system for more than a few orbits by direct numerical integration of the exact equations are prohibitive. ",
        "introduction": "The construction of several new gravitational wave detectors, including the Laser Interferometer Gravitational Wave Observatory (LIGO), TAMA, VIRGO and GEO, may soon make gravitational wave astronomy a reality. The inspiral and coalescence of compact binaries, consisting of neutron stars or black holes, are among the most promising sources for detection by these observatories.  It is expected that neutron star/neutron star binaries will spend approximately 16,000 cycles in the LIGO/VIRGO frequency band, neutron star/black hole binaries about 3,500, and black hole/black hole binaries about 600~\\cite{kip}.  To increase the likelihood of detection, and to extract information from the signal, the binary inspiral has to be modeled theoretically and waveform templates have to be constructed.  The evolution of compact binaries proceeds in different stages.  By far the longest is the initial quasi-equilibrium inspiral stage, during which the compact objects move in nearly circular orbits, while the separation between them slowly decreases as energy is carried away by gravitational radiation.  The quasi-circular orbits become unstable at the innermost stable circular orbit (ISCO), where the inspiral enters a plunge and merger phase.  The merger and coalescence happens on a dynamical timescale, and produces either a black hole or, for binary neutron stars, possibly a larger neutron star, which may collapse to a black hole at a later time.  The final stage of the evolution is the ringdown phase, during which the merged object settles down to equilibrium.  Two distinct approaches have commonly been employed to analyze the adiabatic inspiral phase.  Much progress has been made in post-Newtonian studies of compact binaries (see e.g.~\\cite{djs00} and references therein).  Most of these approaches, however, approximate the compact objects as point sources, which neglects important finite-size effects which may be particularly important for neutron star binaries.  Also, PN expansions may not converge sufficiently rapidly in the strong-field region near the ISCO.  Alternatively, compact binaries can be modeled numerically.  Computational constraints currently limit dynamical evolution calculations to at most a few orbits, so that there is no hope to simulate a complete inspiral.  It is possible, however, to numerically model binaries in the adiabatic inspiral phase in a quasi-equilibrium (QE) approximation.  The QE approximation is based on the assumption that the gravitational radiation reaction time scale is much longer than the orbital timescale, so that the binary can be approximated to be in quasi-equilibrium (and quasi-circular orbit) on an orbital timescale. A similar approximation is routinely being used in stellar evolution calculations.  There, the evolutionary timescale is much longer than the hydrodynamical timescale, so that the star can safely be assumed to be in quasi-equilibrium on a dynamical timescale. Quasi-equilibrium models of compact binaries have been constructed both for neutron stars (see, e.g.,~\\cite{bcsst} for corotating and~\\cite{irrotational} for irrotational binary neutron stars) and for black holes~\\cite{c94}.  Even though individual QE models only represent ``snap-shots'' of binaries at a certain separation, it is possible to construct the complete adiabatic inspiral together with the emitted gravitational wave signal using the following scheme~\\cite{stu}.  For each separation, a QE model can be inserted as a matter source into Einstein's equations (this is the ``hydro-without-hydro'' approach demonstrated in~\\cite{bhs}).  Numerically integrating Einstein's equations for this given matter source will then yield the gravitational wave signal and luminosity of the binary at that separation.  Interpolation between a discrete sample of separations then yields the gravitational wave luminosity as a continuous function of separation.  Combining this with the binary's binding energy as a function of separation, one can construct the inspiral rate at all separations, and hence the separation as a function of time.  Given a suitable parameterization, the entire continuous inspiral wavetrain can then be constructed.  The viability of this approach has recently been demonstrated in~\\cite{dbs} where a prototype calculation was presented for corotating binary neutron stars obeying a polytropic equation of state.  Note that unlike post-Newtonian approaches, the QE scheme uses the fully nonlinear Einstein equations.  The QE approach is also computationally very efficient, since it is possible to produce a very large number of cycles from a small number of QE configurations, each of which needs to be followed for only a couple of orbital periods.  In this paper, we evaluate the QE approach for a model problem in relativistic scalar gravitation, and show that it produces excellent agreement with an exact numerical solution.  While scalar gravity (see, e.g., problem 7.1 in Misner, Thorne and Wheeler~\\cite{mtw}) is not the correct theory of gravitation, it is conceptionally much simpler than general relativity (GR), and shares many of its characteristic features.  This makes scalar gravity a very attractive framework for calibrating the QE approximation scheme, since, unlike in GR, we can directly compare its results with the readily producable exact solution.  In scalar gravity, the gravitational field is described by one scalar function.  Also, scalar gravity generates gravitational radiation even in spherical symmetry, so that the generation and emission of gravitational waves can be studied with a spherically symmetric numerical code, involving only one spatial coordinate.  This enables us to track the effect of radiation reaction exactly over many dynamical timescales.  Moreover, outgoing wave boundary conditions can be imposed correctly at arbitrarily close separations from the sources in spherical symmetry, which eliminates the need for large computational grids.  The theory also admits a local law of energy conservation, while GR only obeys global energy conservation.  In numerical work, such a conservation law provides a strong check on the accuracy of integration.  For all of these reasons, scalar gravitation has been employed successfully to develop many tools of numerical relativity (see, e.g., Shapiro and Teukolsky~\\cite{slst}, hereafter ST, and also~\\cite{slst2}) and we extend that tradition here.  In this paper we study, in the framework of scalar gravitation, the damped oscillations of a spherical matter shell, which we adopt as a simple spherically symmetric analogue to binary inspiral in GR.  A relativistic binary emits gravitational waves, which slowly extracts energy from the binary orbit, so that the inspiral proceeds along a sequence of nearly periodic circular orbits of decreasing separation. Similarly, an oscillating matter shell emits gravitational waves, which slowly extracts energy from the oscillation, so that the damping proceeds along a sequence of nearly periodic oscillations of decreasing amplitude.  Similar to relativistic binaries, where we consider the quasi-circular orbits at a certain separation of the adiabatic inspiral to be in QE, we may also consider the matter shell's quasi-periodic oscillations of a certain amplitude to be in QE.  One difference between the two processes is that the binary inspiral continues until coalescence and merger, while the damping of the matter shell's oscillations will continuously slow down until a true equilibrium state has been reached.  We adopt three distinct approaches to computing the damped oscillation of the matter shell.  We first compute the exact solution, by numerically integrating the exact equations. In our second approach, we adopt a QE approximation by neglecting gravitational waves, and computing QE models of the periodic, undamped oscillations of matter shells for various oscillation amplitudes.  These QE models are then inserted as matter sources into the dynamical equations for the gravitational field.  Integrating the field equations for these QE matter sources yields the gravitational wave form and luminosity for each oscillation amplitude. Combining the wave luminosity with the QE oscillation energy as a function of amplitude yields the amplitude decay rate together with the continuous gravitational wavetrain. This is the analogue to the QE approach to binary inspiral as outlined above.  In a third approach, we use a monopole formula to compute the gravitational waveforms and wave luminosity for QE models and hence to determine the inspiral rate.  This is the equivalent to using the quadrupole formula in GR to compute the inspiral rate and waveforms for binary models.  We compare these three different approaches for three different initial configurations, representing weak, moderately strong and ultra-strong field cases.  We find that for the weak field configuration, all three approaches agree very well.  For the moderately strong field configuration, the QE approach still agrees very well with the exact solution, while the agreement with the monopole result is much worse.  Only for ultra-strong fields, for which the assumptions of QE break down, do we find a disagreement between the QE approach and the exact solution.  However, even in this case, the break-down is gradual and not abrupt.  This is a very encouraging result, and suggests that the QE approach is a very reliable and efficient framework for computing adiabatic binary inspiral up to the ISCO.  The paper is organized as follows.  We summarize the basic equations of scalar gravity in Sec.~II, and outline the QE scheme in the Sec.~\\ref{QES}.  In Sec.~\\ref{numres} we present our numerical results, and compare our three approaches for the three different situations.  We summarize and discuss the implications of our findings in Sec.~\\ref{conc}.  We also include three Appendices.  Appendix A contains a short proof that a minimum in the quasi-equilibrium energy corresponds to a static equilibrium shell solution. Appendix B presents some details of the QE scheme for constructing a continuous wavetrain. Appendix C describes our numerical implementation of the field equation.  Throughout the paper we adopt gravitational units with $G=c=1$. ",
        "conclusions": "\\label{conc}  In this paper, we illustrate our QE approach for calculating gravitational waveforms for a model problem in scalar gravitation. We demonstrate why it is a very viable technique to modeling binary inspiral in GR.  Compact binaries, consisting of neutron stars or black holes, emit gravitational wave and slowly spiral towards each other until they reach the ISCO.  Outside of the ISCO, the inspiral is very slow: the gravitational radiation reaction timescale is much longer than the orbital period.  It is therefore reasonable to assume the binaries to be in QE, and the inspiral to proceed along a sequence of QE configurations.  In a recent paper, Duez, Baumgarte and Shapiro~\\cite{dbs} demonstrated how the inspiral can be modeled by inserting QE binary models as matter sources in Einstein's field equations.  Integrating Einstein's equations then yields the gravitational wave luminosity, from which the inspiral rate and hence the entire gravitational wavetrain can be constructed.  In this paper, we study a simple analogue in scalar gravitation. Scalar gravitation has the advantage that it is conceptionally simpler, and that it admits gravitational radiation even in spherical symmetry, so that we can reproduce all qualitative characteristics of the QE approximation even in a $1+1$ dimensional problem.  We study the radiative damping of an oscillating, spherical matter shell.  In analogy with the orbiting binary in GR, the oscillating shell emits gravitational radiation, and hence looses energy.  While the binary separation decreases in the process, the oscillation amplitude decreases.  Ultimately, the oscillation is completely damped out, and a static equilibrium configuration is reached.   The advantage of scalar gravitation is that it is possible to integrate the exact equations without any approximation.  We compare these exact results with predictions from both our QE approach and the monopole approximation (which is the analogue of the quadrupole approximation in GR).  We find that all three approaches agree very well for weak field solutions, but that the QE approach reproduces the exact solution much better for moderately strong fields.  Only for ultra-strong fields does the QE approach break down.  We conclude that the QE approximation is a very promising approach to computing the adiabatic inspiral of compact binaries."
    },
    "0011/astro-ph0011545_arXiv.txt": {
        "abstract": "{ We consider the effect of vertical outflows on the mean-field dynamo in a thin disk. These outflows could be due to winds or magnetic buoyancy. We analyse both two-dimensional finite-difference numerical solutions of the axisymmetric dynamo equations and a free-decay mode expansion using the thin-disk approximation. Contrary to expectations, a vertical velocity can enhance dynamo action, provided the velocity is not too strong. In the nonlinear regime this can lead to super-exponential growth of the magnetic field. ",
        "introduction": "Magnetic fields play an important role in accretion disks. Inside the disk, for example, they allow the magneto-rotational instability to develop \\citep{BH98}. This instability can lead to the development of turbulence which in turn transports angular momentum radially outwards, thus allowing the matter to accrete. On the other hand, large scale magnetic fields are often considered necessary for launching winds or jets from accretion disks. The origin of such large scale magnetic fields is still unclear: they may be advected from the surrounding medium or be generated by a dynamo inside the disk. Advection appears unlikely as turbulence leads to enhanced viscosity and magnetic diffusivity, so that the two are of the same order \\citep[Prandtl number of order unity;][]{PFL76}. Thus, turbulent magnetic diffusivity can compensate the dragging of the field by viscously induced accretion flow \\citep{vBa89,LPP94,HPB96}. Dynamo action is a plausible mechanism for producing magnetic fields in accretion disks \\citep{Pud81,SL88}. However, dynamo magnetic fields can affect winds from accretion disks \\citep{CPA98,Cam99,BDS00}. On the other hand, the wind can also affect the dynamo. In particular, an enhancement by wind was suggested by \\citet{Bra93} in the context of galactic dynamos. In the same context, \\citet{EGR95} also consider the effects of shear in the wind.  In this paper, we consider the generation of magnetic field by a dynamo acting in a thin accretion disk. We provide a study of the effects of vertical velocities. We show how vertical velocities can enhance the dynamo, allowing a larger growth rate and leading to super-exponential growth of the magnetic field in the nonlinear regime.  The vertical velocities considered can have several origins. They can, for example, be due to a wind emanating from the disk or to magnetic buoyancy inside the disk. Note that we are interested here in magnetic buoyancy as generating vertical outflows. We neglect the associated dynamo effect \\citep{MSS99}.  The paper is organized as follows. In Sect.~2, we describe the two complementary approaches used and specify our choice of parameters. In Sect.~3, we discuss the linear results without vertical outflows. Results with vertical outflows are discussed in Sect.~4 for the linear regime and in Sect.~5 for the nonlinear regime. Sect.~6 presents our conclusions. ",
        "conclusions": "Disk dynamos have been studied intensively over the past 20 years, but only recently the effects of outflows and winds have been included (e.g. \\citealt{Cam99,DBS99,BDS00}). It was immediately clear that these changes in the boundary conditions affect the structure of the solutions. The present work is an attempt to provide a more thorough survey of solutions with wind for different values and sign of the dynamo number. One of the effects that we have highlighted is the enhancement of the kinematic dynamo growth rate for intermediate strengths of the wind. In the nonlinear case this corresponds to a phase shortly before saturation where the growth becomes super-exponential.  In the present work we have deliberately isolated the dynamo problem from the wind problem. In reality the two aspects together constitute a more complicated self-consistent dynamo problem that needs to be looked at in more detail in future. One step in that direction was taken by \\citet{BDS00} where the feedback on the flow from the Lorentz force was taken fully into account. A more formidable task would be to obtain self-consistent dynamo action in three dimensions without assuming an $\\alpha$-effect, thus obtaining dynamo action and outflows in a self-consistent manner from first principles. In the context of accretion disks, this would require simulating turbulent dynamo action, where the turbulence itself is driven by dynamo-generated magnetic fields \\citep{BNS95}. This would be of great interest not only from the astrophysical point of view, but also from that of dynamo theory in general, because, as was recently realized, the catastrophic quenching of the standard $\\alpha\\omega$-dynamo could be alleviated in the presence of open boundary conditions which may substantially improve the working of the dynamo \\citep{BF00}. This is simply because the $\\alpha$-effect generates helical fields, but at the same time magnetic helicity is conserved in a closed domain. Thus, being able to get rid of helical fields through open boundaries may be vital for the dynamo \\citep{BD01}. Furthermore, the production of fields at large scales could be facilitated by the production -- and subsequent loss through boundaries -- of small scale fields with opposite sign of magnetic helicity. This latter aspect could not be addressed with mean-field models and would require a direct simulation of accretion disk turbulence in a global model.  \\acknowledgement{We acknowledge financial support from PPARC (Grant PPA/G/S/1997/00284) and from the Leverhulme Trust (Grant F/125/AL). We thank D. Sokoloff for helpful discussions. The use of the PPARC supported parallel computer GRAND in Leicester and the PPARC funded Compaq MHD Cluster in St Andrews is acknowledged.}"
    },
    "0011/astro-ph0011259_arXiv.txt": {
        "abstract": "indent} \\def\\titlehang{\\hangindent\\abstractindent \\hangafter 0 \\relax} \\newdimen\\secskipamount  \\secskipamount=1pt \\newdimen\\subskipamount  \\subskipamount=1pt \\newdimen\\bottomtol \\bottomtol=0.03\\vsize \\def\\secskip{\\par \\ifdim\\lastskip<\\secskipamount \\removelastskip \\fi \\vskip 0pt plus \\bottomtol \\penalty-250 \\vskip 0pt plus -\\bottomtol \\relax \\vskip\\secskipamount plus3pt minus3pt} \\def\\subskip{\\par \\ifdim\\lastskip<\\subskipamount \\removelastskip \\fi \\vskip 0pt plus 0.5\\bottomtol \\penalty-150 \\vskip 0pt plus -0.5\\bottomtol \\relax \\vskip\\subskipamount plus2pt minus2pt} \\def\\subsubskip{\\par \\ifdim\\lastskip<\\subskipamount \\removelastskip \\fi \\vskip 0pt plus 0.5\\bottomtol \\penalty-150 \\vskip 0pt plus -0.5\\bottomtol \\relax \\vskip\\subskipamount plus2pt minus2pt \\hskip 10pt} \\long\\def\\aaabstract#1{\\centerline{\\null} \\vskip 1.52cm {\\absnarrower \\noindent {\\bf Summary.} #1 \\par} \\oneskip \\oneskip} \\outer\\def\\unnumberedsectionbegin #1 #2\\par {\\secskip \\noindent {{\\bf  #1} \\dotfill #2} \\nobreak \\vskip 1pt \\noindent} \\outer\\def We present observations of the region between $ 5^{\\circ} \\le \\ell \\le 45^{\\circ}$ and $ |b| \\le 3^{\\circ}$, in the OH 1612.231 MHz line, taken from 1993 to 1995 with NRAO's Very Large Array \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. } (VLA). These observations are the last part of a larger survey, covering $ |\\ell| \\le 45^{\\circ}$ and $ |b| \\le 3^{\\circ}$ , with the Australia Telescope Compact Array (ATCA) and the VLA. The region was systematically observed on a 30\\arcmin$\\times$30\\arcmin\\ grid in ($\\ell,b$) and the resulting coverage was 92\\% , with 965 pointings. We found 286 OH--masing objects, 161 of which are new detections and 207 have reliable IRAS point--source identifications. The outflow velocity was determined for 276 sources. A total of 766 sources were detected in the combined ATCA/VLA survey, of which 29 were detected in two regions of the survey.  {\\it The source tables and spectra (Table \\SRCT\\ and \\Fg\\SRCF\\ can be downloaded from http://msowww.anu.edu.au/~msevenst/pubs.html}  In this article we analyse the data statistically and give identifications with known sources where possible. The ``efficiency'' of this VLA survey is 75\\% of that of the ATCA Bulge survey. This efficiency was determined by comparing the detections in the region where the two surveys overlap. The completeness-- and error characteristics are similar, though less homogeneous, except for the much larger errors in the flux densities. The relatively large surface number density found in the northern disk, suggests that we can see the Bar extending to higher longitudes on this side of the \\gc . ",
        "introduction": "We have surveyed a large section of the galactic plane in the OH 1612--MHz satellite line. The aim of the survey was to sample the stellar dynamics in the plane, cutting through the major components of the Galaxy. Very strong radiation is emitted in the observed OH line by, amongst others, OH/IR stars (for a review see Habing 1996). These are particularly suitable tracers of the stellar dynamics for a variety of reasons. Most important is that they are easily observable in any region of the Galaxy, their line--of--sight velocity can be determined very accurately and they represent a large fraction of the stellar population. This article is the sequel to previous articles discussing the Australia Telescope Compact Array (ATCA) ``Bulge'' ($|\\ell|$$<$10\\degr , Sevenster \\etal1997a, Paper I) and ``Disk'' ($\\ell$$<$$-$10\\degr , Sevenster \\etal1997b, Paper II) parts of the survey. Here we discuss the results of the northern plane region, between $ 5^{\\circ} \\le \\ell \\le 45^{\\circ}$ and $ |b| \\le 3^{\\circ}$. This region was observed with the VLA in New Mexico.  \\begin{figure} \\fignam\\COV \\anfig {\\psfig{figure=ms10125f1.ps,width=9cm}} \\vskip -6.5truecm \\caption[]{ The 965 pointings of the VLA survey that were used in the data reduction; no data were obtained for 88 pointings out of 1053. In reality the maps were square (41\\arcmin $\\times$41\\arcmin\\ in $\\alpha,\\delta$) and overlapping. A single missing pointing (of which there are 50, cf.~around $\\ell$=42\\degr ) leaves a hole of $\\sim$10\\arcmin $\\times$10\\arcmin\\ in the coverage, but influences the sensitivity in a larger area. } \\end{figure}  \\begin{figure*} \\fignam\\RFI \\anfig {\\psfig{figure=ms10125f2a.ps,angle=270,height=4.3cm}} {\\psfig{figure=ms10125f2b.ps,angle=270,height=4.3cm}} \\vskip -8.5truecm \\hskip 6truecm{ {\\psfig{figure=ms10125f2c.ps,angle=270,width=12cm}}} \\caption[]{{\\bf a--c.} The two larger panels on the left show the raw spectrum of a primary calibrator (1331+305), taken on 210594, averaged over all baselines (top) and after flagging with WSRFI (bottom). The smaller panels on the right show the single--integration, single--baseline spectra (ie.~visibilities) that were actually flagged by WSRFI. Most of these visibilities were flagged as `G' (\\Sct 3.1), except for the 7th plot, which was a `W'. } \\end{figure*}  We discuss the observations in \\Sct 2 and the data reduction in \\Sct 3. In \\Sct 4 we give the results of this survey and in \\Sct 5 a statistical analysis of the data, that, like Paper II, follows closely the analysis in Paper I. We summarize in \\Sct 6. ",
        "conclusions": "We have given the results of a survey of the region $ 5^{\\circ} \\le \\ell \\le 45^{\\circ}$ and $ |b| \\le 3^{\\circ}$ in the OH 1612.231 MHz maser line. The survey is 99\\% complete for sources brighter than 500 mJy and 90\\% complete for positional offsets from the pointing centres smaller than 13\\arcmin . The absolute flux density limit is $\\sim$120 mJy. We have found 286 compact OH--maser sources, 161 of which are new detections. The sources are mainly OH/IR stars, with a few related sources, like planetary nebulae. The positions are accurate to 2\\arcsec , the velocities to 1 \\kms\\ and the flux densities to 20\\% . For 207 sources, an associated IRAS point source is found. The total number of sources in the combined ATCA/VLA survey is 766.  The VLA survey is similar to the ATCA surveys in terms of global completeness statistics, but considerably less homogeneous in quality. The efficiency is approximately 75\\% that of the ATCA Bulge survey. The biggest differences are the considerably worse accuracy in the flux densities and the very low fraction of single--peaked sources. Signatures of a higher fraction of bar--like stars in this part of the survey than in the ATCA Disk survey are found in the surface number density of detected sources and arguably in the IRAS two--colour diagram."
    },
    "0011/astro-ph0011290_arXiv.txt": {
        "abstract": "\\small{ {\\it Compact symmetric objects} conform a class of sources characterized by high luminosity radio emission located symmetrically on both sides of the active galactic nucleus on linear scales of less than 1 kpc. Given their small size, the hot spots of the jets in CSOs provide a unique laboratory for the study of the physics of relativistic jets and their environment close to the central engine.\\\\ We present a simple model for the hot spots in CSOs assuming synchrotron emission, minimum energy and ram pressure equilibrium with the external medium. Further comparison of our model with observational data allows us to constrain the physical parameters in the hot spots and the jets feeding them, and the density profile of the external medium.} ",
        "introduction": "In the early eighties Phillips and Mutel (1982) first detected compact radio sources with double structure, known as {\\it compact doubles}. Later on, {\\it compact symmetric objects} (CSOs) were defined by Wilkinson et al. (1994) as core-detected sources with double sided emission and total linear size smaller than 1 kpc. They present a typical synchrotron spectrum at high radio frequencies, with a peak due to absorption, around the GHz, i.e., they are also included within a class of objects known as {\\it gigahertz peaked spectrum sources}.  The main controversy about the nature of CSOs has been if they are old, confined sources (van Breugel et al. 1984; Carvalho 1998) or the young precursors of large Faranoff-Riley II (Owsianik et al. 1998; Taylor et al. 2000). The study of this class of compact sources is thus very interesting as a family of objects on their own, and they provide a scenario to probe jet properties close to the active nucleus.  In this contribution, we present results for the physical conditions in the jets and hot spots in a sample of CSOs and summarize the main conclusions of a simple evolution model, trying to discern between the proposed scenarios by comparing the model with observational data. All the quantities have been calculated assuming synchrotron emission, minimum energy and ram pressure equilibrium with the external medium. Calculations are performed assuming the Universe is described by the standard Friedmann-Robertson-Walker model with Hubble constant $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$ and deceleration parameter $q_{0} = 0.5$.  Sources have been selected from the GPS samples of Stanghellini et al. (1998), Snellen et al. (1998), and Peck et al. (2000). We have chosen those sources with double morphology and already classified in the literature as CSOs and also those whose components can be safely interpreted as hot spots even though the central core has not been yet identified. Moreover, sources possibly affected by orientation effects (beaming, spectral distortion), like quasars and core-jet sources, have not been considered. The resulting sample is formed by 20 sources which are listed, along with the data relevant for our study, in Perucho \\& Mart\\'{\\i} (2001). ",
        "conclusions": ""
    },
    "0011/astro-ph0011129_arXiv.txt": {
        "abstract": "I review the statistical techniques needed to extract information about physical parameters of galaxies from their observed spectra. This is important given the sheer size of the next generation of large galaxy redshift surveys. Going to the opposite extreme I review what we can learn about the nature of the primordial density field from observations of high--redshift objects. ",
        "introduction": "Most of the information about the physical properties of galaxies comes from their electromagnetic spectrum. It is therefore of paramount importance to be able to extract as much physical information as possible from it. In principle, it is straightforward to determine physical parameters from an individual galaxy spectrum. The method consists in building synthetic stellar population models which cover a large enough range in the parameter space and then use a merit function (typically a $\\chi^2$) to evaluate which suite of parameters better fits the observed spectrum. There are two obvious limitations of the above method: first, the number of parameters that govern the spectrum of a galaxy may be very large and thus difficult to explore fully. Secondly, in the case of ongoing large redshifts surveys which will provide us with about a million galaxy spectra, it will be computationally very expensive  (and possibly intractable for redshift surveys like the 2dF and SDSS) to apply a plain $\\chi^2$ to each individual spectrum which itself may contain of the order of $10^3$ data points.  The non-obvious route to tackle the problem is to compress the original data set in order to weight more those pixel in the spectrum that carry most information about a given parameter. It is worth reminding that non--optimal data compression is commonly applied to galaxy spectra: photometric filters. Not surprisingly, this empirical data compression is not optimal since it has not been devised to be lossless, i.e. contain all the information for a given parameter. For example, the photometric $B$ filter alone is not optimal to recover the age of a galaxy. On the other hand, more sophisticated and non-empirical methods have been proposed for extracting information from galaxy spectra, some of them as old as the Johnson's filter system itself. Many of these are based on Principal Component Analysis or wavelet decomposition \\cite{Murtagh87,Francis92,Connolly95,Folkes96,Galaz98,Bromley98,Glazebrook98,Singh98,Connolly99,Ronen99,Folkes99}. PCA projects galaxy spectra onto a small number of orthogonal components. The weighting of each component corresponds to it's relative importance in the spectra. However while these components appear to correlate well with physical properties of galaxies, their interpretation is difficult since they do not have known, specific physical properties; they can be amalgams of different properties. To interpret these components, we have to return to model spectra and compare them with the components \\cite{Ronen99}. This is a disadvantage of PCA since one important goal of the analysis is to study the evolution of the physical properties which dramatically affect galaxy spectra, such as the age, metallicity, star formation history or dust content. More sophisticated methods have been recently proposed \\cite{HJL00,SSTL00}. Here I will concentrate in describing the optimal parameter extraction method proposed by \\scite{HJL00}.  The main idea of the method in \\scite{HJL00} is that, in practice, some of the data may tell us very little about the parameters we are trying to estimate, either through being very noisy, or through having no sensitivity to the parameters. So in principle, we may be able to throw some data without loosing much information about the parameters. It is obvious that throwing away data is not the most optimal way. On the other hand, by performing linear combinations of the data we will do better and then we can throw the linear combinations which tell us least. Given a set of data {\\bf x} (in our case the spectrum of a galaxy) which includes a signal part ${\\bf \\mu}$ and noise ${\\bf n}$, i.e. $\\bx = \\bbmu + \\bn$, the idea then is to find a weighting vector ${\\bf b}$ such as $y \\equiv {\\bf b}^{t} {\\bf x}$, it is these {\\it numbers} $y$ which we are after.  In \\scite{HJL00} an optimal and lossless method was found to calculate ${\\bf b}$ for multiple parameters (as is the case with galaxy spectra). Specifically:  \\begin{equation} \\bb_1 = {\\bC^{-1} \\bbmu_{,1}\\over \\sqrt{\\bbmu_{,1}^t \\bC^{-1}\\bbmu_{,1}}} \\label{Evector1} \\end{equation}  and  \\begin{equation} \\bb_m = {\\bC^{-1}\\bbmu_{,m} - \\sum_{q=1}^{m-1}(\\bbmu_{,m}^t \\bb_q)\\bb_q \\over \\sqrt{\\bbmu_{,m} \\bC^{-1} \\bbmu_{,m} - \\sum_{q=1}^{m-1} (\\bbmu_{,m}^t \\bb_q)^2}}. \\label{bbm} \\end{equation}  where a comma denotes the partial derivative with respect to the parameter $m$ and $C$ is the covariance matrix with components $C_{ij}=<n_in_j>$.  The specific steps to build $m$ linear combinations to estimate $m$ parameters are the following:  \\begin{enumerate}  \\item Choose a ``fiducial'' model (a first guess)  \\item Compute the mean spectrum for the $m$ parameters and $m$ partial derivatives with respect to the $m$ parameters ($\\bbmu_{,m}$).  \\item Now compute $m$ eigenvectors $\\bb_{i}$ from Eq.\\ref{Evector1} and \\ref{bbm}.  \\item Finally, compute the $m$ $y_i$ values. This dataset is orthonormal, so the new likelihood is easy to compute (the $y_m$ have mean $\\langle y_m \\rangle = \\bb_m^t \\bbmu$ and unit variance), namely:  \\be \\ln{\\cal L}(\\theta_\\alpha) = {\\rm constant} - \\sum_{m=1}^{M} {(y_m-\\langle y_m\\rangle)^2\\over 2} \\ee  \\end{enumerate}  \\begin{figure} \\centerline{\\psfig{figure=jimfig1.eps,height=6cm,angle=-90}} \\caption{Full likelihood solution using all pixels. There are 6 contours running down from the peak value in steps of 0.5 (in $\\ln{\\cal L}$), and 3 outer contours at $-100$, $-1000$ and $-10000$.  The triangle in the upper-right corner marks the fiducial model which determines the eigenvectors to set the initial weights.} \\label{tth1} \\end{figure} \\begin{figure} \\centerline{\\psfig{figure=jimfig2.eps,height=6cm,angle=-90}} \\caption{Likelihood solution using the compressed data set, i.e. the age datum and the normalization datum.  Contours are as in Fig. \\ref{tth1}.} \\label{tth2} \\end{figure}  This procedure can be applied to derive the metallicities, ages, star formation rates and dust content of galaxy spectra \\cite{RJH00}.  It is very instructive to illustrate the method by trying to recover the age and normalization of single stellar populations (SSP), i.e. the star formation rate is $SFR(t') = A \\delta(t'+t)$ where $\\delta$ is a Dirac delta function. The two parameters to determine are age $t$ and normalization $A$. We built a simulated spectra with Gaussian noise and variance given by the mean, $\\bC = {\\rm diag}(\\bbmu_1, \\ldots)$.  This is appropriate for photon number counts when the number is large. It should be stressed that this is a more severe test of the model than a typical galaxy spectrum, where the noise is likely to be dominated by sources independent of the galaxy, such as CCD read-out noise or sky background counts.  In the latter case, the compression method will do even better than the example here (e.g. \\scite{RJH00}). The simulated galaxy spectrum is one of the galaxy spectra (age 3.95 Gyr, model number 100), and the maximum signal-to-noise per bin is taken to be 2.  Noise is added, approximately photon noise, with a Gaussian distribution with variance equal to the number of photons in each channel.  Hence $\\bC =$ diag($\\bbmu_1,\\bbmu_2,\\ldots)$.  Figure~\\ref{tth1} shows the contours in the likelihood surface using all the points in the spectra. Figure~\\ref{tth2} shows the contours in the likelihood surface using {\\it only two linear combinations: y$_1$ and y$_2$}. As it transpires from the figures, only two numbers suffice to determine two parameters. ",
        "conclusions": ""
    },
    "0011/astro-ph0011403_arXiv.txt": {
        "abstract": "Weak isolated X-ray precursor events before the main Gamma Ray Burst, GRB, and also rare Soft Gamma Repeaters, SGR, events are in disagreement with any Fireball, or Magnetar, scenarios. These models are originated by an unique explosive event leading, by internal-external shock waves, to softer secondary trains following a main gamma signals. Indeed the earliest $GRB980519$, $GRB981226$ events as well as the latest and most distant identified one as GRB000131 are showing rare but well identified and distinct X Ray precursor, occurring tens of seconds or even a minute before the main GRB eruption. These weak X precursors bursts correspond to non-negligible energy powers, up to million Supernova ones. They are rare, about $(3-6)\\%$ of all GRBs, but not unique. Similar huge explosive precursor are in total disagreement with a successive main Fireball GRB outburst. Comparable brief X-ray precursor flashes are found also in rarest and most detailed SGRs events as those observed on 27 and 29 August 1999 from SGR 1900+14. They are inconsistent with a Magnetar Fireball explosion. We interpret them as earlier marginal blazing of outlying X conical Jet tails surrounding a narrower gamma precessing,spinning beamed Jet in blazing mode toward the Earth; later re-crossing and better hitting of the target -the satellite detectors- is source of the main GRB (and SGR) observed structured event. The X Ray precursor existence is an additional remarkable evidence of the Precessing relativistic Jet Nature of both GRBs and SGRs. ",
        "introduction": "Gamma Ray Burst and Soft Gamma Repeaters reached an apparent stage of maturity: tens of GRBs found, finally, an X, optical and (or) radio transient (the after-glow) identification as well as some associated host galaxies at cosmic red-shifts (Bloom et all 2000). New categories of GRBs and SGRs events have been labeled, but even within these wider updated data no conclusive theory or even partial understanding seem to solve the old-standing GRB/SGRs puzzle: the Nature of The GRB-SGR signals. \\\\   On the contrary the wider and wider collection of data are leading to a schizophrenic attitude in the most popular isotropic models, the Largest Cosmic Explosions (Fire-ball, Hypernova, Supra-Nova) with more and more phenomenological descriptions (power laws everywhere) and less and less unifying views. This \"give up\" attitude seem to reflect the surprising never ending morphologies of GRBs. \\\\ We argued on the contrary that GRBs and SGRs find a comprehensive theory within a thin spinning and multi precessing $\\gamma$ Jet, sprayed by a Neutron Star, NS, or Black Hole, BH, (Fargion 1994-1999, Fargion, Salis 1995-1998). For instance the extreme energy released in last GRBs ($\\gg$ $10^{54}$ $erg $), comparable to few solar masses or more, leads to a deep conflict with any isotropic GRB energy (masses and corresponding Schwarchild scale times above milliseconds) and the sharp observed GRBs fine time structures (below a fraction of millisecond). Moreover the energy power spread (from $10^{53}$ $erg s^{-1}$) for most far GRBs versus $10^{46}$  $erg  s^{-1}$ for nearest GRB980425, led most Fireball defenders to neglect, hide or even reject in a very arbitrary way the  nearest and best identified GRB connection to a Supernova, SN, explosion as SN1998bw. Finally the same rarity of GRB-SN detections and the established GRB980425-SN link favors the thin Jet Nature of GRBs.\\\\  Even originally ($1970-80$) unified  GRB/SGR models since last fifteen years are commonly separated by their repeater and spectra differences; however  very recently they openly shared the same spectra, time and flux structures. This favours once again their common Nature. However they are up present times usually described by catastrophic spherical explosions, but, as we shall show in this paper, they should not be. Their different distances, cosmic versus galactic ones, imply different power source Jet, but their morphological similarity strongly suggest an unique process: the blazing of  a spinning and multi-precessing gamma Jet, from either Neutron Star or Black Hole. The $\\gamma$ Jet is born by high GeVs electron pairs Jet which are regenerating, via Inverse Compton Scattering, an inner collimated beamed $\\gamma$ (MeVs) precessing  Jet. The thin jet (an opening  angle inverse of the electron Lorentz factor, a milli-radiant or below), while spinning, is driven by a companion and/or an asymmetric accreting disk in a Quasi Periodic Oscillation (QPO) and in a Keplerian multi-precessing blazing mode: its $\\gamma-X$ ray lighthouse trembling and flashing is the source of the complex and wide structure of observed Gamma Bursts. \\\\ These $\\gamma$ Jets share a peak power of a Supernova ($10^{44} erg s^{-1}$) at their birth (during SN and Neutron Star formations), decaying by power law $\\sim t^{-1}$ $-$ $\\sim t^{-(1.5)}$ to less power-full Jets that converge to present persistent SGRs stages. Indeed these ones  are blazing events from late relic X pulsar observable only at nearer distances. The $\\gamma$ Jet emit in general at $\\sim$ $ 10^{35}$ erg$ s^{-1}$ powers; both of GRB and SGR show an apparent luminosity amplified by the inverse square of the thin from $10^{-3}$ to $10^{-4}$ radiant angle Jet beaming: the corresponding solid angle $\\Omega$ spreads between $10^{-7}$ and $10^{-9}$.\\\\ Optical-Radio After-Glows are not the fading fireball explosion tails often observed in puzzling variable non monotonic decay, but the averaged external Jet tails moving and precessing and geometrically fading away. The rare optical re-brightening (the so called SN bump) observed in few afterglow has been erroneously associated to an underlying isotropic SN flash: it is more probably the late re-crossing of the precessing Jet periphery toward the observer direction.      One of the most recent and convincing evidence against any explosive GRB model, confirming present $\\gamma$ precessing Jet theory, is hidden in the in the recent GRB 000131 data which show an un-explicable, for Fireball model, X precursor signal 7 sec long, just  62 seconds prior to the huge main gamma trigger. How could any GRB source coexist such a powefull precursor?   \\begin{figure}[thb] \\begin{center} \\mbox{\\epsfig{file=fig06.eps, width= 8cm}} \\caption{\\em {Location and Intensity of early $X$ Precursor in GRB000131}} \\end{center} \\end{figure}   \\begin{figure}[thb] \\begin{center} \\mbox{\\epsfig{file=fig01.eps, width= 6cm, bb=10 20 400 200}} \\mbox{\\epsfig{file=fig02.eps, width= 6cm, bb=10 20 400 200}} \\caption{\\em {Fig $2a$ and $2b$: Time evolution and X precursors in GRB $971210$ and GRB $971212$ }} \\end{center} \\end{figure}  \\begin{figure}[thb] \\begin{center} \\mbox{\\epsfig{file=fig03.eps, width= 6cm, bb=10 20 400 200}} \\mbox{\\epsfig{file=fig04.eps, width= 6cm, bb=10 20 400 200}} \\caption{\\em {Fig $3a$ and $3b$: Time evolution and X precursor in GRB $980125$ and GRB $990518$ }} \\end{center} \\end{figure}  \\begin{figure}[thb] \\begin{center} \\mbox{\\epsfig{file=fig05.eps, width= 6cm, bb=10 20 400 200}} \\mbox{\\epsfig{file=fig07b.eps, width= 7cm, bb=10 20 300 200}} \\caption{\\em {Fig $4a$ and $4b$: Time evolution and X precursor in GRB $991216$ and in SGR $1900+14$ on 29 August $1998$ }} \\end{center} \\end{figure} ",
        "conclusions": "GRBs and SGRs  are persistent blazing flashes from light-house thin $\\gamma$ Jet spinning  in multi-precessing (binary, precession, nutation)  mode. These Jets are originated by NSs or BH in binary system or disk powered by infall matter: the Jet is not a single explosive event even in GRB they are powered at maximal output during SN event.The Jet power is comparable at its peak  the $\\gamma$ Jet has a chain of progenitor identities: it is born in most SN and or BH birth and it is very probably originated by very collimated primary muon pairs at GeVs-TeVs energies. These muons could cross the dense target matter around the SN explosions. These muon progenitors might be  themselves secondary relics of pion decays or even by more transparent beamed ultra-high energy neutrino Jet originated (by hadronic and pion showering) near the NS or BH. The relativistic muons  decay in flight in electron pairs is itself source of GeVs relativistic pairs whose Inverse Compton Scattering with nearby thermal photon is the final source of the observed  hard $X$ - $\\gamma$ Jet. The relativistic morphology of the Jet and its multi-precession is the source of the puzzling complex $X$-$\\gamma$ spectra signature of GRBs and SGRs. Its inner internal Jet contain, following the relativistic Inverse Compton Scattering,  hardest and rarest beamed GeVs-MeVs photons (as the rarest EGRET GRB940217 one) but its external Jet cones are dressed by softer and softer photons. This   onion like multi Jets is not totally axis symmetric: it doesn't appear on front as a concentric ring serial; while turning and spraying around it is deformed (often) into an elliptical off-axis concentric rings preceded by the internal Harder center leading to a common Hard to Soft GRBs (and SGRs) train signal. In our present model and simulation this internal effect has been  here neglected without any major consequence. The complex variability of GRBs and SGRs are simulated successfully by the equations and the consequent geometrical beamed Jet blazing leading to the observed $X-\\gamma$ signatures. As shown in fig 5 the slightly different precessing configurations could easely mimic the wide morphology of GRBs as well as the surprising rare X-ray precursor shown in Fig.1-4 above."
    },
    "0011/astro-ph0011153_arXiv.txt": {
        "abstract": "We study H~I Ly$\\alpha$ absorption observed by the {\\em Hubble Space Telescope} toward the nearby binary system $\\alpha$~Cen (G2~V+K0~V) and its distant companion star Proxima~Cen (M5.5~Ve).  Absorption from heliospheric H~I heated by the solar wind/ISM interaction is observed toward both $\\alpha$~Cen and Proxima~Cen.  Absorption from analogous ``astrospheric'' material surrounding the stars is detected toward $\\alpha$~Cen, but not Proxima~Cen.  The nondetection of astrospheric absorption toward Proxima~Cen suggests that the stellar wind of Proxima~Cen must be significantly weaker than that of the $\\alpha$~Cen system.  We use hydrodynamic models of the astrospheres computed assuming different mass-loss rates to predict astrospheric Ly$\\alpha$ absorption for comparison with the observations.  The model that best matches the $\\alpha$~Cen data has a mass-loss rate of $\\dot{M} = 2~\\dot{M}_{\\odot}$, and the models suggest an upper limit of $\\dot{M} \\leq 0.2~\\dot{M}_{\\odot}$ for Proxima~Cen. Finally, we note that the heliospheric absorption observed toward Proxima~Cen in 2000 May is identical to the heliospheric absorption observed toward $\\alpha$~Cen in 1995 May, implying that the structure of the outer heliosphere does not change significantly during the solar activity cycle. ",
        "introduction": "Coronal winds analogous to the solar wind have proven to be very difficult to detect around cool main sequence stars like the Sun.  This is not surprising given the low density of the solar wind, corresponding to a mass-loss rate of only about $\\dot{M}_{\\odot}=2\\times 10^{-14}$ M$_{\\odot}$ yr$^{-1}$, and the fact that the wind is fully ionized.  These properties mean that such winds cannot be detected spectroscopically like the massive winds of hot stars and evolved cool stars.  Ionized winds would be expected to produce radio emission, and searches for this emission have provided upper limits for the mass-loss rates of many solar-like stars.  However, these upper limits are typically 3 or more orders of magnitude higher than the solar mass-loss rate \\citep{ab90,sad93,jl96b,ejg00}.  There have been some claims of high mass-loss rates detected for a few very active stars using observations at millimeter wavelengths and studies of UV absorption features \\citep{djm89,djm92}.  However, these interpretations of the data remain highly controversial \\citep{jl96a,jl96c}.  There are theoretical arguments for large mass-loss rates from active stars \\citep{gdc76,djm92}, but \\citet{jl96c} argue that massive ionized winds would completely absorb the coronal radio emission that is commonly observed from these stars.  Fortunately, a new method for indirectly detecting winds around cool main sequence stars has recently become available, using spectroscopic observations of stellar Ly$\\alpha$ lines made by the {\\em Hubble Space Telescope} (HST).  Models of the interaction between the solar wind and local interstellar medium (LISM) predict that charge exchange processes should create a population of heated neutral hydrogen gas throughout the heliosphere \\citep{vbb95,gpz96,gpz99b,hrm00a}.  This material produces a detectable absorption signature in the Ly$\\alpha$ lines of many nearby stars \\citep{kgg97,vvi99,bew00b}. However, not only is heliospheric absorption detected in the data, but analogous ``astrospheric'' absorption from material surrounding the star is also observed in some cases.  Astrospheric Ly$\\alpha$ absorption has by now been detected for seven nearby coronal stars, although some of the detections are tentative \\citep{bew96,ard97,kgg97,bew98,bew00a}.  This absorption can only be present if a stellar wind is present, and if the star is surrounded by ISM material that is at least partially neutral.  A larger stellar mass loss rate will result in a larger astrosphere and more Ly$\\alpha$ absorption.  \\citet{hrm00b} used this fact to estimate mass-loss rates for two stars ($\\epsilon$~Ind and $\\lambda$~And) based on the amount of observed astrospheric absorption.  In this paper, we perform a similar analysis to estimate the mass-loss rate of the $\\alpha$~Cen binary system (G2~V+K0~V), which has also been observed to have detectable astrospheric Ly$\\alpha$ absorption \\citep{jll96,kgg97}.  In addition, we present new HST Ly$\\alpha$ observations of $\\alpha$~Cen's distant companion star Proxima~Cen (M5.5~Ve), which can be compared directly with the $\\alpha$~Cen data to search for differences in astrospheric absorption that would indicate differences in stellar wind properties.  Unlike the two $\\alpha$~Cen stars, which will share the same astrosphere since their orbit has a semimajor axis of only 24~AU \\citep{dp99}, Proxima~Cen will have an astrosphere all to itself.  Proxima~Cen is about 12,000~AU from $\\alpha$~Cen, based on its $2.2^{\\circ}$ separation on the sky and its closer distance of 1.295~pc compared with 1.347~pc for $\\alpha$~Cen \\citep{macp97}.  Thus, while the LISM and heliospheric absorption should be identical toward $\\alpha$~Cen and Proxima~Cen, the astrospheric absorption should be different. ",
        "conclusions": ""
    },
    "0011/astro-ph0011479_arXiv.txt": {
        "abstract": "Deep submillimetre(submm)-wave surveys made over the last three years using the SCUBA camera at the 15-m James Clerk Maxwell Telescope (JCMT) have revealed a new population of very luminous high-redshift galaxies. The properties of this population, and their contribution to the intensity of the extragalactic background radiation field are described briefly, especially in the context of the SCUBA lens survey.\\footnote{The SCUBA Lens Survey has been carried out since 1997 by Ian Smail, Rob Ivison, Jean-Paul Kneib and the author: a full description of the survey and references to supporting work can be found in the recent catalogue paper \\cite{Smail_cat}} The potential problems caused by source confusion in the large 15-arcsec SCUBA observing beam for the identification and follow up of the results are discussed. The effects of confusion are not important for the study of 850-$\\mu$m SCUBA sources brighter than about 2\\,mJy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/gr-qc0011114_arXiv.txt": {
        "abstract": "s{This talk reviews the constraints imposed by binary-pulsar data on gravity theories, and notably on ``scalar-tensor'' theories which are the most natural alternatives to general relativity. Because neutron stars have a strong gravitational binding energy, binary-pulsar tests are qualitatively different from solar-system experiments: They have the capability of probing models which are indistinguishable from general relativity in weak gravitational field conditions. Besides the two most precise binary-pulsar experiments, in the systems B1913+16 and B1534+12, we also present the results of the various ``null'' tests of general relativity provided by several neutron star--white dwarf binaries, notably those of gravitational radiation damping.}  Binary pulsars allow us to test the strong-field regime of gravity because neutron stars are compact objects: Their mass to radius ratio is of order $Gm/Rc^2 \\sim 0.2$, as compared to $\\sim 10^{-6}$ for the Sun. The best way to illustrate the various tests available is to plot their constraints on a generic class of scalar-tensor theories,\\cite{nonpert} defined by the following action: \\begin{equation} S = \\int d^4 x \\sqrt{-g} \\Bigl( R - 2 g^{\\mu\\nu} \\partial_\\mu \\varphi \\partial_\\nu \\varphi \\Bigr) + S_{\\rm matter}\\Bigl[{\\rm matter}; e^{2\\alpha_0\\varphi + \\beta_0 \\varphi^2}g_{\\mu\\nu}\\Bigr]\\ , \\end{equation} where $\\alpha_0$ and $\\beta_0$ are two constants characterizing the coupling of matter to the scalar field. The origin of Fig.~1 ($\\alpha_0 = \\beta_0 = 0$) corresponds to general relativity; the vertical axis ($\\beta_0 = 0$) to Brans-Dicke theory with a parameter $2\\omega + 3 = 1/\\alpha_0^2$; and the horizontal one ($\\alpha_0 = 0$) to theories which are perturbatively equivalent to general relativity, {\\it i.e.}, strictly indistinguishable from it in the weak-field conditions of the solar system.  Lunar Laser Ranging experiments\\,\\cite{etaLLR} and observations of Mercury perihelion shift% \\ exclude the region lying above the thin solid line of Fig.~1, while Very Long Baseline Interferometry\\,\\cite{gammaVLBI} imposes the limit $|\\alpha_0| < 1.4 \\times 10^{-2}$ materialized by a thin horizontal line. This figure confirms that solar-system experiments do not impose any constraint on $\\beta_0$ if $\\alpha_0$ is small enough.  On the other hand, binary-pulsar data impose\\,\\cite{nonpert} $\\beta_0 > -4.5$ even for a vanishingly small $\\alpha_0$. In terms of the Eddington parameters $\\beta^{\\rm PPN}$ and $\\gamma^{\\rm PPN}$, which are both consistent with 1 in the solar system, this inequality reads $(\\beta^{\\rm PPN}-1) / (\\gamma^{\\rm PPN}-1) < 1.1$. This result comes mainly from the Hulse-Taylor binary pulsar,\\cite{1913} which excludes the region lying above the bold line labeled ``1913+16''.  The next strongest constraint is imposed by neutron star--white dwarf binaries like PSR 0655+64, which excludes the region outside the two solid lines of Fig.~1. Indeed, in scalar-tensor theories, such dissymmetrical systems emit strong dipolar gravitational waves ($\\propto 1/c^3$), generally inconsistent with the small observed value of their $\\dot P_b$ (rate of change of orbital period). The quadrupolar radiation ($\\propto 1/c^5$) predicted by general relativity is consistent with experimental data.  The binary pulsar\\,\\cite{1534} PSR 1534+12 excludes the region lying to the left and top of the dashed line. This is a slightly weaker constraint than with the previous systems, but this test is nevertheless very important, because it does not depend on the radiative structure of the theory.  Scalar-tensor theories also predict a violation of the strong equivalence principle (SEP), which causes a gravitational analogue of the Stark effect on the orbit of dissymmetrical binaries.\\cite{Stark,Wex00} The small observed eccentricities of such systems constrain the theories to lie between the two (approximate) dotted lines.  Similar analyses can also be performed to test local Lorentz invariance of gravity and conservation laws, and notably the ``PPN'' parameters\\,\\cite{alpha1PSR,Wex00} $\\alpha_1$ and\\,\\cite{alpha3,Wex00} $\\alpha_3$.  \\begin{figure} \\begin{center}\\leavevmode\\epsfbox{Fig1.ps}\\end{center} \\caption{Solar-system \\& binary-pulsar constraints on generic scalar-tensor theories. The hatched region is allowed by all the tests. The doubly hatched one is also consistent with VLBI data.$^3$ \\label{fig1}} \\end{figure} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011365_arXiv.txt": {
        "abstract": "We estimate the production of gravitinos during and after the end of a period of warm inflation, a model in which radiation is produced continuously as the field rolls down the potential producing dissipation.  We find that gravitino production is efficient for models in the strong dissipation regime, with the result that standard nucleosynthesis is disrupted unless the magnitude of the inflaton potential is very small. Combining this with the constraint from the thermal production of adiabatic density perturbations we find the dissipation rate must be extraordinarily strong, or that the potential is very flat. ",
        "introduction": "The {\\em warm inflation} scenario \\cite{warm} is an unusual variant on the inflationary cosmology, in which the inflaton has significant interactions during the inflationary epoch leading to a continuous production of radiation. The backreaction of this production on the inflaton field appears as a viscosity, slowing down the scalar field evolution and hence aiding slow-roll inflation \\cite{YML}. In such a scenario, inflation can proceed with potentials steeper than those in standard chaotic scenarios.  Issues concerning the implementation of warm inflation within a realistic particle physics context have yet to be studied to the same depth as the standard inflationary scenario \\cite{LR,LL}. Nevertheless, given that the warm inflationary scenario is very different phenomenologically from the usual picture, it makes good sense to examine the extent to which its phenomenology is consistent with observations. The two main purposes of inflation are to provide a large, nearly homogeneous, patch in the Universe within which structure formation can take place, and to ensure that unwanted relic particles do not spoil the successes of the standard hot big bang cosmology. The first of these has seen a reasonable amount of study \\cite{TB2000}, and so we will consider an example of the latter.  In the context of modern particle physics, the most troublesome relics are the gravitino and the moduli fields \\cite{Sarkar}. We will consider the gravitino, whose existence arises as the supersymmetric partner of the graviton, and whose mass is expected to be order of 1~TeV. It is a cosmological threat because if produced in enough abundance in the early Universe, it is sufficiently long lived to survive until after nucleosynthesis, at which point its decays spoil the element abundances \\cite{Weinberg}. To avoid this, the ratio of gravitino to photon number densities must be below about $10^{-12}$. The gravitino may be produced both by interactions within a thermal bath \\cite{kawa} and by various non-thermal processes \\cite{nontherm}. In conventional inflationary scenarios, the former gives an important upper limit on the reheat temperature, while the latter may constrain many possible physical processes.  In this paper we explore the consequences of gravitino production during and after warm inflation. Warm inflation differs from conventional inflation in that radiation is constantly produced during inflation, and the radiation density decreases monotonically throughout the evolution, with inflation ending when the radiation density overtakes the inflaton energy density. There is therefore continuous gravitino production during inflation, and also no delay in post-inflationary thermal production due to an intervening (p)reheating period. Consequently the gravitino bound is much harder to satisfy. We will show that the abundance of gravitinos produced during inflation is similar to that produced after inflation, and assess the strength of the constraints this imposes on warm inflation model-building. ",
        "conclusions": "A crucial role of inflation is to ensure that unwanted relic particles do not survive with abundances capable of spoiling the standard hot big bang model. We have studied the production of gravitinos during and after a warm inflationary era, and combined the constraint this gives with the requirement that density perturbations of the correct magnitude are generated.  We have found that although there is continuous gravitino production during warm inflation from interactions in the thermal bath, this does not in itself lead to a very significant extra yield of gravitinos over and above that produced at the end of warm inflation. Nevertheless, avoiding overproduction of gravitinos is much more challenging than in conventional inflationary scenarios, because the radiation density is monotonically decreasing throughout the evolution. Satisfying the gravitino bound requires that the potential energy at the end of warm inflation be very small, $V^{1/4} \\lesssim 10^{-9} m_\\Pl$. In order for density perturbations to have the correct magnitude requires a dissipation factor $\\Gamma \\approx 10^{-6} m_\\Pl$, and hence a dimensionless dissipation $r \\gtrsim 10^{12}$ unless the slope of the potential is extremely flat. Either way, it is clear that evading overproduction of gravitinos strongly constrains the warm inflation scenario, requiring dimensionless numbers many orders of magnitude away from unity."
    },
    "0011/astro-ph0011015_arXiv.txt": {
        "abstract": "A simple binary model for Mkn~501 is considered under the assumption that the TeV and X-ray periodicity of $\\sim 23$ days, observed during the flaring state in 1997, may be basically interpreted as a doppler-shifted flux modulation due to the orbital motion of the relativistic jet. For the typical jet properties inferred from emission models, we find an intrinsic orbital period of $(6-14)$ yrs and a centre-of-mass distance of $(2-3.5) \\times 10^{16}$ cm, the binary thus being a very close system. If the binary separation corresponds to that at which gravitational radiation becomes dominant, one may obtain a maximum allowed primary mass of $\\sim 10^8\\, M_{\\odot}$ and a corresponding maximum allowed secondary mass in the range of $\\sim (1-9)\\times 10^7\\, M_{\\odot}$ assuming that gas accretion occurs on around the salpeter time scale. Interestingly such values agree with the black hole masses expected from merger scenarios. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011278_arXiv.txt": {
        "abstract": "We consider models in which the (transverse) photon mass is non-zero during inflation and drops to zero non-adiabatically at the end of inflation. Through this process, vacuum fluctuations of the photon field are converted into physical, long-wavelength modes with high amplitude. The resulting spectrum of the field strength is approximately $B_\\ell\\sim \\ell^{-1}$, where $\\ell$ is the relevant coherence scale. With a reasonable model of field evolution we obtain, on comoving galactic scales, a magnetic field strong enough to seed a dynamo mechanism and generate the observed galactic magnetic fields. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011087_arXiv.txt": {
        "abstract": " ",
        "introduction": "The investigation of stellar populations in  a galaxy using broad-band colours is based on an analysis of the distribution of stars in the colour magnitude diagram (CMD), and  the comparison of the observed CMD with a set of model CMDs. In more distant galaxies, individual stars are  difficult to observe, and hence one must interpret the {\\em integrated\\/} light in terms of the stellar populations and use colours of model populations for comparison. Papers using the last-mentioned approach are becoming more frequent recently. For instance, Kong et al. (2000) used a set of 13 intermediate-band colours to get information on the distribution of age, extinction and metallicity in M81, assuming a population of stars of equal age in each pixel. Bell et al (1999) and Abraham et al (1999) analysed the pixel-by-pixel colour distribution to get the overall star formation history, comparing the colours with theoretical predictions for populations with various time-scales $\\tau$ of the star forming process.  In contrast, Notni and Bronkalla (1984) derived the population distribution in the inner regions of M82 assumimg {\\em two distinct}\\/ populations  (each with small $\\tau$) in each pixel in their analysis of photographic data. Similarly, the integrated colour for a sample of dwarf galaxies was interpreted by Almoznino and Brosch (1998) as a combination of a young (3-13Myr) and an old (1-2 Gyr) stellar populations. They get their conclusion by calculating some combinations for the integrated light of two populations with different ratio. We start,  the other way round, from the observed light in every point in the galaxy and try to decompose  this light into its contributions from the two stellar populations by solving the equations which describe the combined colours. The method enables us  to estimate the age and the extinction of the population I and the contribution of both populations I and II to  the observed light.  In the present paper, the mathematical  presentation of the method is given in section 2. A short error analysis in the determination of age and relative contribution of different populations is discussed in section 3. In section 4, an application of the method using CCD observations of the irregular galaxy NGC 3077 is presented. The conclusion will be given in section 5. The observations used in this work and their reduction are presented in our preceeding paper in this issue (Abdel-Hamid \\& Notni 2000, HN1 hereafter). ",
        "conclusions": "We  present a simple method for deducing the relative contributions of the stellar populations of galaxies by an analysis of their colours. We solve the 'mixing equations' (by iteration) for the case of two populations of markedly different age in NGC 3077.  For NGC 3077, the main results of this paper can be summarized as follows:\\\\ - We derive the spatial distribution of both populations together with an estimate of  their  age and an extinction  estimate for population I. \\\\ - The young population is concentrated in the central region of NGC 3077. Its contribution to  the observed light decreases going out from the centre.\\\\ - The youngest structures, aged 10-20 Myr, concentrate along a  ringlike structure,  associated with dust. The ring breaks up into discrete star-forming clumps.\\\\ - An age gradient for the young population in the central region of NGC 3077 is found,  ages range from 4 Myr in the centre to 100 Myr at a projected radial distance of 1 kpc. Such age variation reflects the occurence of  different events of star formation, probably starting from the beginning of the encounter with M81 (a few hundred million years) until now.\\\\ - The prominentest star cluster in the central region is found to have the lowest age, about 4 Myr.  The abovementioned results encourage us to use our method to study the stellar populations in faint galaxies using only the familiar broad band colours."
    },
    "0011/astro-ph0011564_arXiv.txt": {
        "abstract": "We model Galactic populations of accreting and cooling isolated neutron stars in the attempt to explore their link with a new class of dim soft X-ray sources revealed by ROSAT. For accretors we follow the magneto-rotational and dynamical evolution in the Galactic potential and a realistic large scale distribution of the interstellar medium is used. Under standard assumptions old neutron stars enter the accretor stage only if their magnetic field exceeds $\\approx 10^{11}$--$10^{12}$ G. We predict about 1 source per square degree for fluxes $\\approx 10^{-15}$--$10^{-16}$ erg cm$^{-2}$s$^{-1}$ in the energy range 0.5-2 keV.  Cooling neutron stars are explored within a simpler model of local sources, including however interstellar absorption. They are found to be significantly less abundant at low fluxes, $<0.1$ sources per square degree, but dominate over accretors at higher flux levels ($\\approx 10^{-12}$--$10^{-11}$ erg cm$^{-2}$s$^{-1}$). We suggest that the faint sources observed by ROSAT may be young cooling neutron stars with typical age  $\\la 10^6$ yrs, if the total number of young neutron stars in the Solar proximity is $\\sim 10$ times higher than inferred from radiopulsars statistics. ",
        "introduction": "Seven soft sources, {\\it The Magnificent Seven}, have been found in ROSAT fields which are most probably associated with isolated radio-quiet NSs. A summary of their main observational properties is reported in the table  (see e.g. Ne\\\"uhauser \\& Tr\\\"umper \\cite{nt99}, NT99; Motch \\cite{mo2000}; Treves et al. \\cite{t2000}, T2000, and references therein; see also Burwitz et al. at this conference).  \\begin{table} \\label{table1} \\caption{Properties of ROSAT Isolated NS Candidates} \\bigskip \\begin{tabular}{lccccc} Source& PSPC & $T_{bb}$ & $N_H$ & $\\log f_X/f_V$ & Period \\\\ & count$\\, \\rm{s}^{-1}$ & eV & $10^{20}\\, {\\rm cm}^{-2}$ &  & s\\\\  MS 0317.7-6647 & 0.03 & 200 & 40 & $> 1.8$ & -- \\\\ RX J0420.0-5022 & 0.11 & 57 & 1.7 & $>3.3$ & 22.7 \\\\ RX J0720.4-3125 & 1.69 & 79 & 1.3 & 5.3 & 8.37 \\\\ RX J0806.4-4132 & 0.38 & 78 & 2.5 & $> 3.4$ & -- \\\\ RBS1223 & 0.29 & 118 & $\\sim 1$ & $> 4.1$ &   5.2 \\tablenotemark{a} \\tablenotetext{a}{Hambaryan et al. (2001)} \\\\ RBS1556 & 0.88 & 100 & $< 1$ &$ > 3.5$ & -- \\\\ RX J185635-3754 & 3.64 & 57 & 2 & 4.9 & --  \\\\  \\end{tabular} \\end{table}  Present X-ray and optical data however do not allow an unambiguous identification of the physical mechanism responsible for their emission. These sources can be powered either by accretion of the interstellar gas onto old ($\\approx 10^{10}$ yr) NSs or by the release of internal energy in relatively young ($\\approx 10^6$ yr) cooling NSs (T2000 for a recent review). The ROSAT candidates, although relatively bright (up to $\\approx 1 \\ {\\rm count\\,s}^{-1}$), are intrinsically dim and their inferred luminosity ($L\\approx 10^{31} \\ {\\rm erg\\, s}^{-1}$) is near to that expected from either a close-by cooling NS or from an accreting NS among the most luminous. Their X-ray spectrum is soft and thermal, again as predicted for both accretors and coolers (T2000). Up to now only two optical counterparts have been identified: RXJ 1856-37, Walter \\& Matthews (\\cite{wm97}), for which a distance estimate of $\\sim 60$ pc has been very recently obtained,  and RXJ 0720-31, Kulkarni \\& Van Kerkwick (\\cite{kk98}). In both cases an optical excess over the low-frequency tail of the black body X-ray spectrum has been reported.  A statistical approach, based on the comparison of the predicted and observed source counts may provide useful informations on the nature of these objects. Previous studies derived the $\\log N$ -- $\\log S$ distribution of accretors (Treves \\& Colpi \\cite{tc91}; Madau \\& Blaes \\cite{mb94}; Manning et al. \\cite{mann96}) assuming  a NSs velocity distribution rich in slow stars ($v\\la 100 \\ {\\rm km\\, s}^{-1}$). Recent measurements of pulsar velocities (e.g. Lyne \\& Lorimer \\cite{ll94}; Hansen \\& Phinney \\cite{hp97}) and upper limits on the observed number of accretors in ROSAT surveys (Danner \\cite{dan98}) point, however, to a larger NS mean velocity (T2000). Ne\\\"uhauser \\& Tr\\\"umper (NT99) compared the number count distribution of the ROSAT isolated NS  candidates  with those of accretors and coolers. Here we address these issues in  greater detail, in the light of the latest contributions to the modeling of the evolution of Galactic NSs (Popov et al. \\cite{p2000a}, P2000a). A more comprehensive discussion can be found in Popov et al. (\\cite{p2000b}, P2000b). ",
        "conclusions": "The task of explaining the observed properties of the seven ROSAT isolated NS candidates within a single model based on standard assumptions is not easy. The ROSAT sources are in fact: i) relatively bright, $\\ga 0.1$ count s$^{-1}$; ii) close, $N_H\\sim 10^{20} {\\rm cm}^{-2}$; iii) soft, $T_{eff}\\sim 50 - 100$ eV; iv) slowly rotating, for RX J0720, RX J0420 and RBS1223 the periods are about 5--20 s.  The accretion stage  can be reached only if NS magnetic field is $\\ga 10^{11}$--$10^{12}$ G. For polar cap accretion the X-ray spectrum is then relatively hard with a typical temperature around 300-400 eV. For them interstellar absorption is not very significant, and we predict about 1 source per square degree for fluxes about $10^{-15}$--$10^{-16}$ erg cm$^{-2}$ s$^{-1}$ for energy range 0.5-2 keV. While accretors may indeed represent the bulk of a still undetected low-luminosty population of X-ray sources, it is difficult to reconcile them with the relatively bright ROSAT sources, as shown in the figure. On the contrary, the number distribution of coolers is in agreement with that of the seven isolated NS sources discovered so far, only if the total number of young close-by NSs is a factor $\\sim 10$ larger than what implied by radiopulsar statistics. Taken at face value, this implies that the majority of NSs do not experience an active radiopulsar phase, a major implication which is also supported by other evidences (see e.g. Gotthelf \\& Vasisht \\cite{gv2000}). The interpretation of the observed periods of some of the Magnificent Seven remains problematic and seems to require special conditions on the evolution of the magnetic field and spin. \\bigskip  \\noindent {\\bf Acknowledgements} The work of SBP and MEP was supported through the grant RFBR 00-02-17164. Support from the European Commission under contract ERBFMRXCT98-0195 and from MURST under contract COFIN98021541 is acknowledged."
    },
    "0011/astro-ph0011422_arXiv.txt": {
        "abstract": "The Galactic Center contains the most massive young clusters in the Galaxy and serves as the closest example of a massive starburst region. Our recent results suggest that the Galactic Center environment produces massive clusters with relatively flat initial mass functions, as might be expected on theoretical grounds. I will discuss these recent results, along with evidence for star formation in the immediate vicinity of the super massive black hole at the Galactic Center. The results of this work might be useful in extrapolating to other galactic centers with similar conditions, as well as other starburst regions. ",
        "introduction": "The Galactic Center region spans the central few hundred parsecs of the Galaxy and represents about 0.04\\% of the volume in the Galactic disk. Although small in size, the region contains 10\\% of the molecular material and young stars in the Galaxy, as evidenced by molecular line and radio maps (\\cite{mor96}). The molecular clouds in the center tend to have high densities that are greater than 10$^4$~cm$^{-3}$; more rarefied clouds are destroyed by the strong tidal field produced by the large stellar mass peaking in the central parsec. In addition, the infrared energy distributions produced by the clouds indicate very high temperatures, T$_{\\rm cloud}$~=~70~K; and molecular line observations indicate extraordinary turbulent velocities (FWHM) of about 25~km~s$^{-1}$, even on scales as small as 0.3~pc. In accordance with approximate pressure balance, the magnetic field strengths appear to also be high, on the order of a milliGauss. \\cite{mor96} provides an excellent review of the GC environment.  Despite this tumultuous environment, the GC contains stars which produce L~=~10$^9$~\\Lsun, and L$_{\\rm Ly-cont}$~=~(1$-$3)(10$^{52}$)~s$^{-1}$, values dominated by older stars in the case of the former and younger stars in the case of the latter. Given this fertile breeding ground for young stars, we embarked on a program to identify all the sites of recent star formation in the GC and measure the initial mass function for particular sites, where possible (\\cite{fig95a}; \\cite{fig95b}; \\cite{fig96}; \\cite{fig98}; \\cite{fig99a}; \\cite{fig99b}). This paper gives a summary of the results obtained so far, with a particular emphasis on the most recent result, a measurement of the initial mass function (IMF) in the Galactic Center. ",
        "conclusions": "The first direct determination of the IMF has been made in the Galactic Center for the Arches cluster, and it appears to be significantly flatter than the Salpeter value. This result implies that the IMF law is not universal and may be affected by environmental conditions, although a similar result was found for NGC3603 \\cite{eis98}. We are continuing attempts to minimize possible errors in this determination. For instance, Keck spectroscopy has recently been obtained to verify the mass-magnitude relations. Finally, higher resolution imaging is needed to extend the IMF down to lower masses and to determine the lower-mass cutoffs."
    },
    "0011/astro-ph0011108_arXiv.txt": {
        "abstract": "s{We investigate the behavior of the magnetic pressure, $b^2$, in fully turbulent MHD flows in ``1+2/3'' dimensions by means of its effect on the probability density function (PDF) of the density field. We start by reviewing our previous results for general polytropic flows, according to which the value of the polytropic exponent determines the functional shape of the PDF. A lognormal density PDF appears in the isothermal ($\\gamma=1$) case, but a power-law tail at either large or small densities appears for large Mach numbers when $\\gamma >1$ and $\\gamma < 1$, respectively. In the isothermal magnetic case, the relevant parameter is the field fluctuation amplitude, $\\dbb$. A lognormal PDF still appears for small field fluctuations (generally the case for {\\it large mean fields}), but a significant low-density excess appears at large fluctuation amplitudes ({\\it weak mean fields}), similar to the behavior at $\\gamma > 1$ of polytropic flows. We interpret these results in terms of simple nonlinear MHD waves, for which the magnetic pressure behaves linearly with the density in the case of the slow mode, and quadratically in the case of the fast wave. Finally, we discuss some implications of these results, in particular the fact that the effect of the magnetic field in modifying the PDF is strongest when the mean field is weak.} ",
        "introduction": "\\label{sec:intro}  A fundamental feature of compressible turbulence is the formation of density fluctuations, a property  of central interest in astrophysics, as density excesses (``clouds'') in the turbulent interstellar medium (ISM) exhibit numerous statistical properties over a range of sizes\\cite{larson81,blitz93} whose origin is not yet well understood. Although a full understanding of star formation requires knowledge of the full (multiple-point) statistics in order to determine mean densities as a function of region size, a first step towards this goal is the description and physical understanding of the one-point statistics, or {\\it probability density function} (``PDF'') of the mass density field.  Previous studies have shown that the density PDF in turbulent compressible flows has a lognormal form in the isothermal case,\\cite{VS94,PNJ97,OGS99,OSG00} but exhibits a power-law tail at densities larger (smaller) than the mean for flows with polytropic exponents smaller (larger) than unity.\\cite{scalo_etal98,PVS98,NP99} It should be noted that some of those works referred to purely hydrodynamic flows,\\cite{VS94,PVS98} while the rest referred to magneto-hydrodynamic (MHD) flows.  In particular, ref.\\ [\\cite{PVS98}] presented a heuristic model for the development of the lognormal and power-law PDFs, based essentially on the behavior of the speed of sound with density in those types of flows. In this sense, the density PDF is a diagnostic for the dependence of the pressure with density. However, such model applied only to non-magnetic flows, while the real ISM is most likely magnetized to a significant extent. In this paper we present preliminary results on the density PDF of turbulent magnetized flows in ``1+2/3'' dimensions, as a first attempt to characterize the behavior of magnetic pressure with density in the fully turbulent regime. Previous works have focused on the pressure produced by weakly nonlinear Alfv\\'en waves\\cite{MZ95,Spangler}, but here we consider fully turbulent regimes with arbitrarily large magnetic fluctuation amplitudes. ",
        "conclusions": "\\label{sec:concl}  \\subsection{Summary} \\label{sec:summary}  In this paper we have reviewed our previous results\\cite{PVS98} on the development of the density PDF as a consequence of the effective (polytropic) equation of state, and applied them to an understanding of the PDF in the magnetic case.  In the context of the model presented in ref. [\\cite{PVS98}], isothermal non-magnetic flows have lognormal mass density PDFs while in the non-isothermal cases, a power-law tail develops at high Mach number, at either $\\rho > \\langle \\rho \\rangle$ or $\\rho < \\langle \\rho \\rangle$ depending on whether $\\gamma <1$ or $\\gamma > 1$, respectively.  In the magnetic case, we considered only isothermal cases and a propagation parallel to the ambient magnetic field, but reported that a deviation from the lognormal PDF occurs {\\it when the magnetic fluctuations are large}, a case expected when the mean field is small. We interpreted this effect in terms of so-called ``simple'', finite-amplitude nonlinear MHD waves. Indeed, for the slow mode of these waves, the magnetic pressure $b^2$ behaves linearly with the density $\\rho$, while for the fast mode, the magnetic pressure behaves quadratically with $\\rho$. Together with the fact that for the slow wave the density is weighted by a small factor, this suggests that the production of density fluctuations is dominated by the slow waves at small field deviations (i.e., $\\Bo$ large) giving a lognormal PDF again, while for weak mean fields (large field deviations) the quadratic density dependence of $b^2$ produces an excess at small densities in the PDF, corresponding to the $\\gamma > 1$ case of the polytopic description. A more detailed study, including in particular the case of non-parallel propagation is in progress.  \\subsection{Astrophysical Implications} \\label{sec:astro_impl}  The results of this paper have a number of implications in the astrophysical context. First, we have seen that the main parameter determining the effect of the magnetic field appears to be the magnetic fluctuation amplitude, $\\dbb$, rather than $\\beta$ alone, which is the parameter most frequently used to characterize MHD flows. In other words, it is important to have a knowledge of the relative importance of the turbulent fluctuations (as given by the Alfv\\'enic Mach number) in addition to simply the ratio of thermal to magnetic pressures (as given by $\\beta$).  Second, as far as the density PDF is concerned, the effect of the magnetic field is most notorious when the field fluctuations are large (generally, when the mean field is weak). This suggests that the limit of vanishing field does not approach the non-magnetic case, and in this sense the latter is singular. This may imply that it is inadequate to model the ISM as a non-magnetic flow even if the degree of magnetization is very low, as has started to be recently contended at the level of molecular clouds (see, e.g., ref.\\ [\\cite{PN99}])."
    },
    "0011/astro-ph0011458_arXiv.txt": {
        "abstract": "Galaxy morphological and spectroscopic types should be nearly independent of apparent magnitude in a local, magnitude-limited sample. Recent luminosity function surveys based on morphological classification of galaxies are substantially more successful at passing this test than surveys based on spectroscopic classifications.  Among spectroscopic classifiers, those defined by small aperture fibers (ESP, LCRS) show far stronger systematic classification biases than those defined by  large apertures (APM). This effect can be easily explained as an \\emph{aperture bias}, whereby galaxies with globally late-type spectra are assigned earlier spectral types which depend on the redshift and luminosity of the galaxy.   The effect is demonstrated by extracting successively larger aperture spectra from long-slit spectroscopy of nearby galaxies. This systematic classification bias is generic, and the 2dFGRS  and SDSS surveys will show similar systematic biases in their spectral types.  We suggest several methods for correcting the problem or avoiding it altogether.  If not corrected, this aperture bias can mimic galaxy evolutionary effects and distort estimates of the luminosity function. ",
        "introduction": "The classification of galaxies into different types is more than mere phenomenology: it is a method to separate galaxies in a well-defined manner in order to study their common properties locally and their evolution with redshift.  The applications of galaxy classification are quite diverse, since significant differences in measurable physical properties can be attributed to various galaxy types: luminosity functions for distinct galaxy types evolve differently with redshift (\\cite{lilly95}); galaxies differ in their clustering properties through the  morphology-density relation (\\cite{dressler80}) and their bias with respect to the  underlying mass distribution (\\cite{hubble36}; \\cite{oemler74}; \\cite{blanton99});  and the luminosity evolution of early-type (\\cite{dokkum96}, \\cite{kochanek00b}) and late-type (\\cite{vogt96}) galaxies at fixed mass is also very different for $0<z<1$ due to their different stellar populations.  All these properties are relevant for estimates of the local galaxy power spectrum (\\cite{kolatt97}) and for estimates of the cosmological model using gravitational lenses (e.g. \\cite{kochanek96}).  Galaxies can be classified based on morphology, spectra, colors, and surface photometry to produce a one-dimensional sequence.  While spectral and morphological methods, for example, yield similar classification information (\\cite{morgan57}), these two methods are not identical (\\cite{connolly95}).  Each method emphasizes a different combination of physical properties such as current star formation and AGN activity (emission lines, color), past star formation (absorption lines, color), disk instabilities (spiral structure), and merger histories (bulge-to-disk ratios).  The continuous one-dimensional sequence is then broken into discrete types at somewhat arbitrary boundaries. Finally, galaxies are assigned specific types, usually without uncertainty estimates, and the results are analyzed.  Thus, while the classification of galaxies is necessary to studies of the structure, formation, and evolution of galaxies, the differences and uncertainties in classification methods create many problems of interpretation.   Morphological classification has a long history dating back to Hubble (1936\\markcite{hubble36}), and has several ``picture books'' which are typically referenced to create the operational definition of the morphological types---the ``Hubble Atlas of Galaxies'' (\\cite{hubbleatlas}) and the Revised Shapley-Ames Catalog (\\cite{rsa}).  The method has had a recent renaissance with its application to galaxies at cosmological distances using the high spatial resolution and sensitivity of the Hubble Space Telescope (e.g., field galaxies in the HDF, \\cite{vandenbergh96}; clusters galaxies, \\cite{dressler97}). Spectroscopic classification of galaxies is a more recent field (\\cite{connolly95}) which offers the promise of objective decomposition of spectra into various galactic components (nucleus, bulge, and disk), but may suffer from the reduced spatial information typically available for each galaxy.  While it might be simple to equate morphological and spectroscopic classifications, their precise relationship to each other will also be some function of the observational limitations, such as spatial resolution, rest-frame wavelength of the filter, spatial extent used to extract the spectra, etc.  It is still an open question as to which classification method---morphological, spectroscopic, bulge-to-disk, color, and so on---has superior precision, reliability, and physical insight.  In Kochanek \\etal\\ (2000\\markcite{kochanek00a}) we derived the first local infrared galaxy luminosity function large enough to be compared to optical galaxy luminosity functions.  As part of our analysis we morphologically classified the galaxies because comparisons of the infrared and optical results have to be done by galaxy type due to the large differences in the $(B-K)$ or $(R-K)$ colors of galaxies of different morphological or spectral types.  We used morphological classification because it was relatively easy to do for our bright, local galaxy sample and because we relied so heavily on archival redshifts that we lacked the data necessary to attempt spectral classification.  We could compare our results to eight relatively recent derivations of luminosity functions divided by galaxy type. The APM survey has derived results using both morphological (\\cite{loveday92}) and spectral classification (\\cite{loveday99}). The CfA (\\cite{marzke94b}) and SSRS2 (\\cite{marzke98}) surveys used morphological types.   The LCRS (\\cite{lin96}), ESP (\\cite{zucca97}), and APM (\\cite{loveday99}) surveys used emission line equivalent widths to define types, and the LCRS (\\cite{bromley98a}) and 2dFGRS (\\cite{folkes99}; \\cite{slonim00}) used global analyses of the spectra related to principal component analysis to define types.  The differences in type definitions, magnitude systems, and survey volumes make this comparison difficult, but comparison of all the surveys suggests a qualitative difference in the shapes of the luminosity functions derived from  morphological and spectral classification. Luminosity functions are usually parameterized by the Schechter (1976\\markcite{schechter76}) form, $dn/dL = (n_*/L_*)(L/L_*)^\\alpha \\exp(-L/L_*)$ with a comoving density $n_*$, break luminosity $L_*$, and a faint-end slope $\\alpha$. The differences in luminosity functions derived by morphological and spectroscopic classification are shown in Figure~1. The surveys based on spectroscopic classification show less agreement among them than those based on morphological classification. The surveys using spectral typing also tend to find shallower faint-end slopes for early-type galaxies and steeper slopes for late-type galaxies---as well as larger differences between the two slopes---than do the surveys using morphological typing. These differences for normal galaxies must arise from one of two reasons:  either the differences point to intrinsic differences between the methods and hence intrinsic differences among galaxy properties (morphological versus spectroscopic), or one of the methods is systematically biased.  \\placefigure{fig-lf-morph-spec}  \\begin{figure}[t] \\figurenum{1} \\epsscale{0.5} \\plottwo{f1a.eps}{f1b.eps} \\caption{ Luminosity function shapes for morphologically [left] and spectrally [right] classified galaxy samples.  The Schechter $\\alpha$ parameter (1976) is shown for each survey as a function of galaxy type ordered from early to late type in going from left to right.  If fraction $f_i$ of the galaxies in the survey are of type $i$, then the galaxy fraction assigned to the type is $x_i = \\sum_{j<i} f_j + f_i/2$ which is the sum of the number of galaxies in the earlier bins plus half the number in the current bin.  The APM survey classified by morphological type has 21\\% of  its galaxies unclassified (which have been added as a gap between the early  and late type results), and we used the type fractions of the SSRS2 survey for  the very similar CfA survey since Marzke \\etal\\ (1994b) % did not report the numbers of galaxies in the survey for each type. The studies based on morphological classification show better consistency in their derived luminosity functions than studies based on spectral classification. \\label{fig-lf-morph-spec} } \\end{figure}  Suspicion about galaxy classification usually focuses on morphological classification because of its dependence on biological neural networks and its association with old-fashioned, uncomputerized astronomy.  While many astronomers may feel that the definitions of morphological galaxy types are similar to the famous, if nebulous, Supreme Court definition of pornography (``I know it when I see it,'' \\cite{stewart64}), the studies by Naim \\etal\\ (1995a\\markcite{naim95a},b\\markcite{naim95b}) demonstrated that the results are consistent between different implementations of biological neural networks and can be automated and reproduced using silicon neural networks.  Successful morphological classification does require high dynamic range and resolution images which largely confines its use to the local universe (with ground-based imaging) or intermediate redshifts (with the \\emph{Hubble Space Telescope}).  Morphological classification remains labor intensive and difficult to apply on a very large scale, which has driven most recent surveys to switch to spectral classification.  Although the methods differ in their details, the basic assumption is that the \\emph{global} spectra of galaxies form a continuous sequence between spectra dominated by absorption and emission lines and that this sequence is related to, but not necessarily identical to, the classical morphological sequences (see \\cite{connolly95}; \\cite{zaritsky95}; \\cite{bromley98a}; \\cite{folkes99}; \\cite{slonim00}). The primary advantages of spectral typing are that it is easily automated and that the spectroscopic data required for it are obtained in the course of the redshift survey.  Since, as Figure 1 illustrates for the faint-end slope, the parameters of the various luminosity functions are mutually inconsistent, we started searching for tests and cross calibrations which could be used to reconcile or check the results.  In Kochanek \\etal\\ (2000a\\markcite{kochanek00a}) we used the 2MASS survey to compare directly the magnitude systems in the different surveys. Here we apply a simple test for the self-consistency of luminosity functions by type---namely that the type fractions should be nearly independent of the apparent magnitude---to check whether systematic errors in defining galaxy types can be responsible for the differences in the shapes of luminosity functions. This test is described and applied in \\S2 to the samples for which we could obtain the necessary data. The test is passed by the local, morphologically typed samples and failed by the spectrally typed samples based on small aperture fiber optic spectra. The origin of the problem is described in \\S3 as a probable consequence of the small spectroscopic apertures. The future of using spectra to assign galaxy types is summarized in \\S4. ",
        "conclusions": ""
    },
    "0011/astro-ph0011491_arXiv.txt": {
        "abstract": "The Edinburgh-Cape Blue Object Survey is an ongoing project to identify and analyse a large sample of hot stars selected initially on the basis of photographic colours (down to a magnitude limit $B \\sim 18.0$) over the entire high-Galactic-latitude southern sky, then studied with broadband $UBV$ photometry and medium-resolution spectroscopy.  Due to unavoidable errors in the initial candidate selection, stars that are likely metal-deficient dwarfs and giants of the halo and thick-disk populations are inadvertently included, yet are of interest in their own right.  In this paper we discuss a total of 206 candidate metal-deficient dwarfs, subgiants, giants, and horizontal-branch stars with photoelectric colours redder than $(B-V)_0 = 0.3 $, and with available spectroscopy.  Radial velocities, accurate to $\\sim 10-15$ \\kms , are presented for all of these stars.  Spectroscopic metallicity estimates for these stars are obtained using a recently re-calibrated relation between Ca {\\sc II} K-line strength and $(B-V)_0$ colour.  The identification of metal-poor stars from this colour-selection technique is remarkably efficient, and competitive with previous survey methods.  An additional sample of 186 EC stars with photoelectric colours in the range $-0.4 \\le (B-V)_0 < 0.3$, comprised primarily of field horizontal-branch stars and other, higher-gravity, A- and B-type stars, is also analysed.  Estimates of the physical parameters T$_{\\rm eff}$, log g, and \\feh\\ are obtained for cooler members of this subsample, and a number of candidate RR Lyrae variables are identified. ",
        "introduction": "Deep, and in particular, {\\it complete} surveys of stars in the halo of the Galaxy always turn up objects of astrophysical interest outside of the original targets of the survey.  The Edinburgh-Cape Blue Object Survey (hereafter, EC survey) is no different in this respect (see Stobie et al. 1997 for a detailed description).  Briefly, the EC survey is intended to discover large numbers of stars in the southern sky with colours normally associated with the blue stellar populations of the Galaxy, such as white dwarfs and subdwarf O- and B-type stars, in a similar fashion to the Palomar-Green survey in the northern hemisphere (Green, Schmidt, \\& Liebert 1986).  The initial colour-selection criterion for EC stars makes use of a $B$,$U-B$ polygon in the (photographic) magnitude-colour diagram (see Fig. 2 of Stobie et al. 1997).  The attempt was made, on a plate-by-plate basis, to set the red boundaries of this polygon in such a way as to exclude most of the F- and G-type stars, while retaining most of the bluer objects.  The shape of the boundary differs from plate to plate, and naturally becomes somewhat less clear at fainter magnitudes, due to the increasing errors in the photographic photometry.  For the EC survey, a contamination level of $\\sim $20 per cent in the original selection was deemed acceptable.  During the course of spectroscopic and photometric follow-up observations of the EC stars brighter than $B \\sim 16.5$, it has become clear that many of these `contaminant' stars are in fact of great interest as probes of the metal-poor stellar populations of the Galaxy.  The present paper explores the nature of these stars.  In this paper we report radial velocities, stellar classifications, and spectroscopic abundance estimates, expressed in terms of \\feh\\ , for 321 main-sequence turnoff (TO), subgiant (SG) , giant (G), and field horizontal-branch (FHB) or other A-type stars identified in the process of the EC survey follow-up campaign.  Radial velocities (only) are reported for 73 stars which are likely B-type horizontal-branch (HBB) stars.  We note that the present sample contains 108 stars with radial velocity and abundance determinations in the Galactic anti-rotation direction at lower Galactic latitudes ($|b| \\le 45^\\circ $); these data provide valuable information concerning the rotation of the halo and thick-disk components of the Galaxy (see, e.g., Chiba \\& Beers 2000). ",
        "conclusions": ""
    },
    "0011/astro-ph0011202_arXiv.txt": {
        "abstract": "I review the general aspects of cosmological parameter estimation from observations of the cosmic microwave background (CMB) temperature anisotropies in the framework of inflationary adiabatic models. The most recent CMB datasets are starting to give good constraints on the relevant parameters of inflationary adiabatic models. They point toward a model consistent with the basic predictions of inflation: a nearly flat universe, with a nearly scale invariant spectrum of primordial fluctuation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011034_arXiv.txt": {
        "abstract": "Strongly correlated {\\sl non-intrinsic variability} between 5 and 8.5\\,GHz has been observed in one of the lensed images of the gravitational lens B1600+434.  These non-intrinsic (i.e. `external') variations are interpreted as {\\sl radio-microlensing} of relativistic $\\mu$as-scale jet components in the source at a redshift of $z$=1.59 by massive compact objects in the halo of the edge-on disk lens galaxy at $z$=0.41. We shortly summarize these observations and discuss several new observational and theoretical programs to investigate this new phenomenon in more detail. ",
        "introduction": "Multi-frequency observations of the CLASS gravitational lens (GL) B1600+434 with the VLA and WSRT radio telescopes at 1.4, 5 and 8.5\\,GHz has unambiguously shown {\\sl non-intrinsic variability} in the lensed image that passes through the dark-matter halo of the edge-on disk lens galaxy (Koopmans \\& de Bruyn 2000; Koopmans et al. 2000a). Based on the amplitude--timescale and frequency dependence of these non-intrinsic variations {\\sl and} the difference in variability between the two lensed images, separated by only 1\\farcs4, the {\\sl non-intrinsic variability} is interpreted as {\\sl radio-microlensing} of relativistic $\\mu$as-scale jet components in the source by massive compact objects in the halo of the edge-on disk lens galaxy. The alternatives, i.e.  scintillation and extreme scattering, have considerable difficulties in explaining the time-scale and frequency dependence of the observations (e.g. Koopmans \\& de Bruyn 2000; Koopmans et al. 2000a). An example of these variations at 5\\,GHz is shown in Fig.1. The strong non-intrinsic event between days 70--100 is interpreted as a {\\sl radio-microlensing caustic crossing}. This type of event is extremely difficult to explain in terms of scattering (ISS) by the ionized interstellar medium (e.g. Koopmans et al. 2000b). ",
        "conclusions": ""
    },
    "0011/astro-ph0011020_arXiv.txt": {
        "abstract": "I review the method of \\index{Doppler tomography}Doppler tomography which translates binary-star line profiles taken at a series of orbital phases into a distribution of emission over the binary. I begin with a discussion of the basic principles behind \\index{Doppler tomography}Doppler tomography, including a comparison of the relative merits of maximum \\index{entropy}entropy regularisation versus \\index{filtering}filtered \\index{back projection}back-projection for implementing the inversion.  Following this I discuss the issue of noise in Doppler images and possible methods for coping with it. Then I move on to look at the results of Doppler Tomography applied to \\index{cataclysmic variable}cataclysmic variable stars. Outstanding successes to date are the discovery of two-arm \\index{spiral arms}spiral shocks in \\index{cataclysmic variable}cataclysmic variable accretion discs and the probing of the stream/magnetospheric interaction in magnetic cataclysmic variable stars. \\index{Doppler tomography}Doppler tomography has also told us much about the stream/disc interaction in non-magnetic systems and the \\index{irradiation}irradiation of the \\index{secondary star}secondary star in all systems. The latter indirectly reveals such effects as \\index{shadowing}shadowing by the accretion disc or stream. I discuss all of these and finish with some musings on possible future directions for the method. At the end I include a tabulation of \\index{Doppler map}Doppler maps published in refereed journals. ",
        "introduction": "Many rapidly rotating single and binary stars change little during the course of a single rotation or orbit. The spots on single stars can persist for many days, while \\index{cataclysmic variable}cataclysmic variable stars may stay in \\index{outburst}outburst for over 100 orbits and in \\index{quiescence}quiescence for ten times longer still. However, for the observer, orbital rotation can cause considerable \\index{variability}variability both in flux and spectra. This arises from a combination of changes in aspect angle and \\index{visibility}visibility, caused by geometrical effects, and the rotation of all velocity vectors with the binary orbit. These effects are a blessing and a curse: without them we would know considerably less than we do about such stars, however the complex \\index{variability}variability can be hard to interpret.  The method of \\index{Doppler tomography}Doppler tomography was developed to unravel the emission line variations of \\index{cataclysmic variable}cataclysmic variable stars (CVs) [48].  CVs are short period binary stars, with orbital periods typically between $1.5$ and $10$ hours, which are beautifully set up to allow us to study accretion. The stellar components of the binary, a \\index{white dwarf}white dwarf and a low-mass main-sequence star, are faint, and their semi-detached configuration means that the geometry is entirely specified by the \\index{mass ratio}mass ratio and orbital \\index{inclination}inclination alone. Unfortunately, CVs are far too small to be resolved directly -- they typically subtend $< 10^{-4}$ seconds of arc at Earth -- and we can learn nothing of their structure from direct imaging. Instead we must turn to more indirect methods.  Two key methods are ``eclipse mapping'', introduced by Horne [39] and reviewed in this volume by Baptista, and ``\\index{Doppler tomography}Doppler tomography'', the subject of this chapter. Eclipse mapping relies on the geometrical information contained in eclipse light curves; Doppler Tomography uses the velocity information contained in Doppler-shifted light curves.  In this paper I detail the principles behind \\index{Doppler tomography}Doppler tomography and hope to give the reader a full picture of what is now a widely-applied tool in the field. I follow the fundamentals with a short section on accounting for stochastic noise, which has usually been ignored to date, before finally moving onto a survey of results.  Although I will attempt to be as self-contained as possible, the subject has become too large to cover every application of \\index{Doppler tomography}Doppler tomography and so I instead focus upon \\index{cataclysmic variable}cataclysmic variable stars. I start this endevour with a potted history of the development of \\index{Doppler tomography}Doppler tomography. ",
        "conclusions": "I have reviewed the principles and practice of the analysis method of \\index{Doppler tomography}Doppler tomography which helps the interpretation of the complex line profile variations from close binary stars. The key discoveries from the application of this technique are the \\index{spiral arms}spiral shocks in \\index{outburst}outbursting \\index{cataclysmic variable!dwarf nova}dwarf novae and the stream/\\index{magnetic field}magnetic field interaction in the polar class of \\index{cataclysmic variable}cataclysmic variable stars, but \\index{Doppler tomography}Doppler tomography has taught us much about the stream/disk interaction and \\index{irradiation}irradiation of the \\index{\\index{secondary star}secondary star}donor star too.  For the future, efforts need to be made to acquire multi-epoch datasets for tomography as these are wholly lacking at present. Following that higher spectral and temporal resolution data are needed to exploit tomography to its limit."
    },
    "0011/astro-ph0011216_arXiv.txt": {
        "abstract": "The XMM-OM instrument extends the spectral coverage of the XMM-Newton observatory into the ultraviolet and optical range. It provides imaging and time-resolved data on targets simultaneously with observations in the EPIC and RGS. It also has the ability to track stars in its field of view, thus providing an improved post-facto aspect solution for the spacecraft. An overview of the XMM-OM and its operation is given, together with current information on the performance of the instrument. ",
        "introduction": "The Optical/UV Monitor Telescope (XMM-OM) is a standalone instrument that is mounted on the mirror support platform of XMM-Newton (Jansen et~al.~2001) alongside the X-ray mirror modules. It provides coverage between 170~nm and 650~nm of the central 17 arc minute square region of the X-ray field of view (FOV), permitting routine multiwavelength observations of XMM targets simultaneously in the X-ray and ultraviolet/optical bands. Because of the low sky background in space, XMM-OM is able to achieve impressive imaging sensitivity compared to a similar instrument on the ground, and can detect a $B=23.5$ magnitude A-type star in a 1000~s integration in ``white'' light (6 sigma). It is equipped with a set of broadband filters for colour discrimination. The instrument also has grisms for low-resolution spectroscopy, and an image expander (Magnifier) for improved spatial resolution of sources.  Fast timing data can be obtained on sources of interest simultaneously with image data over a larger field.  In the following sections we give an overview of the instrument followed by an account of its operation in orbit and the instrument characteristics. ",
        "conclusions": "The first stage of commissioning and calibrating XMM-OM has been completed. The instrument is fulfilling its role of extending the spectral coverage of XMM-Newton into the ultraviolet and optical band, allowing routine observations of targets simultaneously with EPIC and RGS. Specifically, the instrument has successfully been demonstrated to provide wide field simultaneous imaging with the X-ray camera, simultaneous timing studies, and boresight information to arcsecond accuracy. A number of results illustrating the scientific potential of XMM-OM are contained within this volume."
    },
    "0011/astro-ph0011485_arXiv.txt": {
        "abstract": "The ``current wisdom'' regarding the structuring of the X-ray emitting corona in active stars (i.e.\\ a corona dominated by extended coronal structures) is briefly reviewed, followed by a review of a new approach to flare analysis and the analysis of a significant number of newly observed and previously published large flares, all leading to a much more compactly structured corona. Recent observations showing the polar location of the flaring plasma are then discussed, showing how the current evidence points toward a (flaring) corona composed of rather low-lying polar structures, also in agreement with some recent radio VLBI observational results and with starspot Doppler images. The resulting picture is significantly different from the solar case. ",
        "introduction": "Active late-type stars have a coronal X-ray luminosity of up to $\\ge 10\\,000$ times higher than the Sun; it is thus natural to ask whether the structuring of such a corona can be simply understood as a simple extension of the solar case or not. Under structuring here we mean the size, location and type of the coronal structures responsible for the bulk of the observed emission measure at X-ray wavelengths (i.e.\\ plasma at MK temperatures). For the Sun, high-resolution X-ray images of the corona have allowed a detailed study of its spatial structuring, showing a corona structured in magnetic loops confined to intermediate latitudes, with ample coronal holes specially at the solar poles. The typical maximum height of the corona above the photosphere is well below the solar radius.  Can this corona be scaled to one $10\\,000$ times more luminous, and if so, how?   Stars are unresolved objects and therefore only indirect evidence can be obtained about the spatial size, location and structuring of their corona. Three basic types of observations can be used to derive structural information about the corona, namely 1) the analysis of flare decay, 2) the modulation of light due either to mutual eclipses in a binary system or to self-eclipse due to stellar rotation in a single star, and 3) the determination of plasma densities from high-resolution spectra.  Here we briefly review the above approaches, showing that in many cases the results have been interpreted, in the past, as giving evidence for a large, rather diffuse corona. We then proceed to discuss recent developments in the methodology as well as in the observations, and discuss how these new development convincingly point toward a picture of a rather compact corona, located on (or near) the poles of active stars. ",
        "conclusions": "The large body of observational evidence presented here (and accumulated mostly in the last two years) shows that the active (flaring) component of the corona of active stars is confined in structures which are significantly smaller than the star itself, and which are likely located at high stellar latitude. This is in contrast with the picture of very long, near-equatorial coronal structures (even with interconnecting loops in the case of active binaries, e.g.\\ Uchida \\& Sakurai 1983) which has in the past been proposed to describe active coronae. Independent evidence for material at significant distance from the stars, and in the inter-binary region comes from Doppler mapping done e.g.\\ with IUE observations (Pagano et al.\\ 2000, on AR Lac), using chromospheric lines. Such evidence is not in contrast with the picture presented here, given than the two concern material at very different temperatures, and thus likely not co-located: the flaring coronal plasma discussed here is mostly at temperatures of several MK to several tens of MK, while the chromospheric material whose presence is deduced through Doppler analysis (using e.g.\\ Mg\\,{\\sc ii} lines) is at temperatures $T \\le 10^5 K$. Similar considerations apply to the evidence for large structures orbiting the young star AB Dor obtained by Collier Cameron et al.\\ (1990): also in this case, the material detected is at much cooler temperatures than the coronal material responsible for the X-ray (flaring) emission, and thus unlikely to be spatially coincident with it.  The polar location of the flaring plasma is in stark contrast with the solar case, where active regions are confined to low and intermediate latitudes. This (together with the evidence from the radio observations of a largely dipolar field) points to a magnetic field structure, in the case of active stars, significantly different from the solar one, and thus also possibly to a different dynamo. Therefore, for the very active stars, the solar analogy may have to be considered with some caution, and the coronal structures responsible for the bulk of the (flaring) X-ray emission may indeed look different than in the Sun.  The paradigm of a compact (and polarly located) corona in active stars is also supported by the EUVE observations of contact (W UMa-type) binaries, for which high density are deduced from spectroscopy (and thus small volumes) for the emitting plasma, and the lack of strong orbital modulation (as well as the period difference with respect to the orbital motion, Brickhouse \\& Dupree 1998 on 44 Boo) also point to a polar compact corona. The upcoming flow of grating spectral observations of active stars from Chandra and XMM-Newton will likely put to test the picture presented here: spectroscopic observations of the decay of an intense flare, once achieved, would allow to independently measure the density and compare it with the one deduced from the hydrodynamic method used here. The ``holy grail'' for the field of coronal structuring would however be a grating observation of an event similar to the SAX Algol one."
    },
    "0011/astro-ph0011399_arXiv.txt": {
        "abstract": "We discuss a class of low--luminosity E/S0 galaxies which have both HI disks and (in contrast to more luminous E/S0s with HI) ongoing star formation. We suggest that such objects are common, but that only a few are known at present because optical magnitude--limited galaxy catalogues are biased against them. The HI Parkes All--Sky Survey (HIPASS) should eventually detect many more. We suggest that `boxy' and `disky' ellipticals are distinct not only in their structure and kinematics, but in their star--formation history. ",
        "introduction": "Twenty years ago, it might have been surprising to see elliptical galaxies featured at a conference on ``Galaxy Disks and Disk Galaxies''. Now, however, it is generally accepted that many elliptical galaxies have stellar and/or gaseous disks.  Intermediate--luminosity ellipticals often show `disky' isophotes and/or rapid stellar rotation suggesting the presence of an inner stellar disk (e.g. Bender et al.\\ 1988; Rix \\& White 1990). Extended disks of gas and dust are often seen around giant ellipticals (e.g. Knapp et al.\\ 1985; Sadler \\& Gerhard 1986), and are usually ascribed to a recent accretion event or merger. Small central disks of ionized gas are also common in elliptical and S0 galaxies (e.g. Phillips et al.\\ 1986; Buson et al.\\ 1993).  The existence of these disks has implications for our ideas about how early--type galaxies form and evolve.  There is no clear photometric boundary between the E and S0 classes (e.g. van den Bergh 1989), and it has been suggested that `disky' ellipticals form a continuous photometric sequence with S0s and spiral bulges and may have a different formation mechanism from the (generally more massive) `boxy' ellipticals (Capaccioli et al.\\ 1990; Kormendy \\& Bender 1996).  For the past few years, we (and others) have been working to obtain high--quality HI images for a larger sample of early--type galaxies in order to learn more about the relationship between gas content, galaxy evolution, environment and the triggering of an active nucleus. Jacqueline van Gorkom (this meeting) has reviewed what we now know about the HI structure and kinematics in early--type galaxies.  In this paper, we focus on two topics --- star formation in the HI disks of low luminosity ellipticals, and the likely impact of the HI Parkes All--Sky Survey (HIPASS) on HI studies of early--type galaxies. In the discussion which follows, we assume H$_\\circ$=50 km s$^{-1}$ Mpc$^{-1}$, and that a `low--luminosity' galaxy has M$_{\\rm B} > -20$\\ mag.  \\vspace*{-0.24cm} ",
        "conclusions": "Luminous (`boxy') and low--luminosity (`disky') ellipticals appear to differ not only in their structure and kinematics, but in their HI content and star--formation history.  HI gas disks are much more common in low--luminosity ellipticals than in high--luminosity ones, and the HI disks in low--luminosity ellipticals commonly support a modest level of ongoing massive star formation, which is not seen in the giant ellipticals. This implies that low--luminosity ellipticals form much more slowly than giant ellipticals, and in many cases are still forming stars today.  \\vspace*{-0.1cm}"
    },
    "0011/astro-ph0011350_arXiv.txt": {
        "abstract": "We  report  the  results  of  spectroscopic   observations  of  8  southern polar-ring galaxies (PRGs), in the wavelength range 5\\,900--7\\,300~\\AA.  We find that 5 out of 8 galaxies  contain  LINERs or Sy  nuclei.  Taking  into consideration all PRGs with available spectral data, we estimate that about half of all PRGs and PRG  candidates  have either LINER or Seyfert  nuclei. The observed widths of the [N\\,{\\sc ii}]$\\lambda$6583 line in the nuclei of early-type  PRGs follow the  linewidth--absolute  luminosity  relation  for normal E/S0 galaxies.  We found that one of the observed  galaxies -- ESO~576-G69 -- is a new  kinematically-confirmed  polar-ring  galaxy with a spiral host. ",
        "introduction": "Polar-ring  galaxies  (PRGs) are rare  relics of galaxy  interactions.  The Polar-Ring  Galaxy  Catalogue  (PRC,  Whitmore et al.  1990)  contains only about 100 PRGs and PRG candidates in the northern and southern hemispheres. At present, at least half of all  northern  PRGs and  candidates  have been investigated,  southern objects have been studied less extensively  (except for a few well-known examples, e.g.  NGC~4650A).  The main aim of the present  work is to collect new  observational  data on southern  hemisphere  PRGs in order to enlarge  the  available  information about this  specific  type of  extragalactic  objects.  We discuss  nuclear spectra of 8 PRGs and candidates of the PRC.  Throughout this work the value H$_0$=75 km $\\rm s^{-1}Mpc^{-1}$ is adopted.  \\section[]{Observations and data reduction}  The   observations   were  performed   with  the   1.6m-telescope   at  the Observat\\'orio  do Pico dos Dias (OPD, Brazil) in May 1997, equipped with a Cassegrain   spectrograph  and  CCD-05-20-0-202   detector  $\\#48$  (thick, front-illuminated  and  coated  for the  UV)  with  770$\\times$1152  square pixels,  22.5\\,$\\mu$m  each,  6.6$e^-$  readout  noise,  and 3.3  $e^{-}\\rm ADU^{-1}$ gain.  The grating of 900 lines $\\rm mm^{-1}$ was centered at 665 nm, with a dispersion of 1.15 \\AA \\,$\\rm  pixel^{-1}$ and resolution of 1.1 pixels   (FWHM).  The  slit  size  was  3   arcsec$\\times$52   arcsec.  The configuration yielded a scale factor of 0.235 arcsec $\\rm pixel^{-1}$ and a spectral  coverage of 5\\,900~{\\AA} --  7\\,300~{\\AA}.  The seeing during the observations was $<$\\,2 arcsec.  The spectra were taken with the slit along the  major  axes of the  central  galaxies  (except  for  AM~1837-631).  In general  three 20 minute  exposures  were taken and coadded of each object, except for  AM~1934--563 of which only one 20 min spectrum was obtained.  A log of the observations is given in Table~\\ref{logobs}.  The reductions were carried out with standard techniques  using\\,{\\sc IRAF} and {\\sc ESO-MIDAS}  packages.  This includes bias subtraction,  flat-field correction,   cosmic  rays   removal,  sky   subtraction   and   wavelength calibration.  \\setcounter{table}{0} \\begin{table} \\caption[]{Log of observations} \\label{logobs} \\begin{tabular}{llll} \\\\ \\hline Galaxy &   PRC          & Date & Exposure    \\\\ &   number       &      &  (min)      \\\\ \\hline ESO~503-G17  & B-12 & 10 May 1997 & 20+20+20         \\\\ Abell~1631-14& B-13 & 10 May 1997 & 20+20+20         \\\\ NGC~5122     & B-16 & 06 May 1997 & 20+20+20         \\\\ AM~1934-563  & B-18 & 06 May 1997 & 20              \\\\ ESO~500-G41  & C-33 & 11 May 1997 & 20+20+20        \\\\ ESO~576-G69  & C-46 & 11 May 1997 & 20+20+20        \\\\ ESO~232-G4   & C-52 & 10 May 1997 & 20+20+20         \\\\ AM~1837-631  & D-29 & 11 May 1997 & 20+20+20        \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "Our results  provide new data on nuclear  properties of PRGs.  A first look at  Table~\\ref{charsum}  shows that active nuclei are  overpresented in the sample, where they represent $\\sim$60 per cent (5/8).  Because of the small sample size, this  conclusion is not  statistically  significant.  However, considering  all 16 objects from  Reshetnikov \\& Combes  \\shortcite{rc}  as well as the data on NGC~2685  \\cite{wef},  NGC~660  \\cite{van}, and IC~1689 \\cite{ht}, the sample of PRGs and candidates with nuclear spectra increases to 27  galaxies.  Of these at least 14 (52 per cent) have  LINER or Seyfert nuclei.  Considered  separately,  and  according  to the  original  papers, Seyfert  nuclei  are found in 3--6  objects  (the  first  number  refers to confident   classifications   and  the  second   one   includes   uncertain classifications) or 12.5--25 per cent,  respectively, and LINERs in a total of 8--11  objects  (counted  as above) or 33--41 per cent.  Therefore,  the fraction of active nuclei is high among PRGs and candidates.  PRGs hosting AGNs are  predominantly  S0 galaxies.  It would be interesting to compare  the number of  S0$_{\\rm(Sey)}$  and  S0$_{\\rm(LINERS)}$  to the number of ordinary S0.  This requires good morphological classifications of our PRGs and candidates, but such comparison is presently  impossible since the   available   data  is  neither  good  nor  abundant   enough  for  the classification.  Moreover, most PRG candidates are very peculiar and faint. More detailed and precise statistics can only be done when we have a sample of PRGs at least twice as large as presently  available  together with good optical images.  PRGs  are  very   heterogeneous   objects  regarding  both  morphology  and environment.  However,  they have one  particular  feature in  common:  the existence of two large-scale strongly inclined kinematic  subsystems.  Such complicated  internal kinematics are considered usually as a consequence of relatively  long-lasting galaxy interactions,  accompanied by mass transfer from one  galaxy  to  another  (ranging  from  gas  accretion  to  complete merging).  One can speculate that such  interactions are favourable for the formation of non-thermal  nuclear activity.  Due to the limited size of our sample this conclusion must, however, remain  tentative for the time being."
    },
    "0011/hep-ph0011386_arXiv.txt": {
        "abstract": "We study the lensing of neutrinos by astrophysical objects. At the difference of photons, neutrinos can cross a stellar core; as a result the lens quality improves. While Uranians alone would benefit from this effect in the Sun, similar effects could be considered for binary systems. ",
        "introduction": "Gravitational deflection of massless particles close to a massive object has been extendedly studied in the case of photons and can be detected either through multiple images or by signal enhancement \\cite{weinberg,gl}.  Here due to the extreme difficulty to detect neutrinos and the comparatively poor angular resolution of \"neutrino telescopes\", the only thing we can hope for is a signal intensification.  While stars are notorously bad lenses for photons, we will show that the situation can be very different for neutrinos: as they can cross even a stellar core, a real focusing of a neutrino beam can be achieved. ",
        "conclusions": "With an approach based on general relativity equations we have studied with some generality neutrino lensing through different models of astrophysical objects. The main difference with photons rests in the possibility for medium energy neutrinos to cross even the stars (despite a cut-off above a few 100~GeV in that case). This has the important result that the quality of the lens is greatly improved, as in that case we get a nearly constant focal length in the central region of the star. This would have been very promising, had the focal length of the center of the Sun happened to coincide with the radius of the Earth orbit; unfortunately this is far from being the case, and the focusing occurs closer to Uranus. For what concerns the Lorentzian case, we notice that the result we obtained is significantly different from the one usually quoted for photons passing through an isothermal potential.  We have considered other astrophysical objects than simple stars, and various geometrical configurations, in order to identify the cases most favourable for lensing. We have taken into account the finite size of the source and of the detector and shown how they can influence the lensing characteristics.  The case of a binary system seems to be quite promising, provided small and energetic sources exist; and galaxy lenses could also provide sizeable enhancements of the signal."
    },
    "0011/astro-ph0011436_arXiv.txt": {
        "abstract": "We have searched for discrete refractive scattering events (including effects due to possible non-multiple diffractive scattering) at meter wavelengths in the direction of two close by pulsars B0950+08 and B1929+10, where we looked for spectral signatures associated with the multiple imaging of pulsars due to scattering in the interstellar medium.  We do not find any signatures of such events in the direction of either source over a spectral periodicity range of 50 KHz to 5 MHz.  Our analysis puts strong upper limits on the column density contrast associated with a range of spatial scales of the interstellar electron density irregularities along these lines of sight. ",
        "introduction": "Radio signals, during their passage through the interstellar medium, are scattered due to the irregularities in the density of free electrons (Scheuer 1968). Most often, signals from distant sources undergo strong and multiple scattering. This leads to many observable effects such as the apparent angular broadening of the source, temporal broadening of the pulse profile, diffractive and refractive scintillations. The nature of interstellar scattering at this limit and the effect on relevant observables have been studied in detail by many early workers (Lee \\& Jokipii 1975; 1976; Blandford \\& Narayan 1985; Blandford, Narayan \\& Romani 1984; Cordes 1986). Refractive alterations (by irregularities larger than the relevant Fresnel scale) of diffraction patterns sometimes lead to multiple images with angular separations larger than the diffraction broadening of the source (Cordes \\& Wolsczan 1986; Cordes et al. 1986; Goodman \\& Narayan 1989a,b).  If the refractive alterations become significant, then multiple imaging manifests itself in the form of periodic spectral \\& temporal modulation of intensity with periods smaller than the decorrelation bandwidth \\& decorrelation time, respectively. The probablity of multiple imaging occuring is low for a Kolmogorov spectrum of density distribution, and increases for steeper spectra as they imply dominant refractice effects (see Cordes et al. 1986 for details). However, for nearby sources, refraction angles may become comparable with the diffraction broadening even for relatively less steep density spectra.  It is also likely that the diffractive scattering of the signals from some of the close-by sources is less `multiple' in nature (with a relatively small number of speckles), mimicing in some sense the multiple imaging situation, where, with an unscattered version of the signal, only a few delayed versions interfere.  Also, this may improve the detectability of a distinct signature of {\\it single} or {\\it discrete refractive scattering events}, where we receive, only a few discrete (diffractive) bundles of rays that correspond to different refracted (and correspondingly) delayed versions of the signal.  In such a case, it appears possible to probe the properties (such as the size and density contrast) of the discrete density-irregularities responsible for the refractive/diffractive scattering, if the associated time delays can be measured. ",
        "conclusions": "In our observations, we have looked for fine periodic spectral structure across a bandwidth not very much wider than the decorrelation bandwidth.  Also, the data span used corresponds to an interval smaller than the decorrelation time for diffractive scintillations of both pulsars. Thus our analysis uses data over at most a couple of diffractive scintles, and is therefore less affected by any possible decoherence due to variation of the diffractive phase from scintle to scintle.  Considering the possibility that the spectral modulation (which we have searched for) could drift across the spectrum with time (although the drift is expected to be small over our data span), we have also examined the average of ACF amplitudes alone (instead of its vector equivalent) so as to avoid possible `dephasing' of any otherwise detectable ACF feature.  Again, no significant feature was apparent.  We interprete our {\\it null} result as indicating an absense of strong refraction gradients across the scattering screens. Our results also suggest that the diffractive scattering is rather `multiple', i.e. the diffraction image consists of a large number of `speckles' or sub-images.  While the former indication can be understood simply as the result of a non-steep spectrum of density irregularities with an `inner scale' much smaller than the Fresnel scale, we attempt to quantify our conclusion using the following simple-minded picture.  Crossing of refracted ray bundles is a necessary condition for multiple imaging (see Cordes et al 1986 for more details). This condition is not met if the differential refraction across a given transverse scale is less than its angular size as seen from the observer's location.  Let us consider a refracting screen mid-way along a pulsar sight-line, and assume that two wavefronts, one of which is refracted significantly more (by, say, an angle $\\theta$) than the other, describe the refractive effect of the scattering screen over a transverse separation `$a$'.  For a difference $\\Delta DM$ in the electron column density sampled by the two paths, the two wavefronts will cross if $\\Delta DM$ $\\ge$ $(8\\pi/Lr_e) (a/\\lambda)^2$, where $r_e$ is the classical electron radius, $L$ is the distance to the pulsar, and $\\lambda$ is the wavelength of observation. For the situation described here, it is easy to show that the geometric delay difference will be equal to the differencial dispersion delay corresponding to $\\Delta DM$.  Figure 3 shows the above mentioned limits on column density contrast $\\Delta DM$ (in units of pc cm$^{-3}$) associated with large-scale irregularity as a function of transverse scale `$a$' (in AU) for the two pulsars, B0950+08 (solid line) and B1929+10 (`dash' line). The distances to these two pulsars are assumed to be 120 pc and 170 pc, respectively. The vertical axis (right side) indicates the associated relative delay (including the dispersion contribution) for comparison with our measurements.  The horizontal axis (top side) is also marked with relative epochs (for the line of sight to traverse a mid-way transverse distance of $a$) corresponding to a proper motion of 100 km/s. It would be instructive to compare these limits on the structure function of dispersion measure when suitable direct measurements of changes in DM in the two directions may become available in future. It is also worth mentioning that the characteristic delay associated with diffractive scattering of the two cases (B0950+08 \\& B1929+10) are 4.7 \\& 0.08 $\\mu$sec, respectively, and would translate to scattering disk size of 0.8 \\& 0.12 AU, respectively. The Fresnel scale, for comparison, is $\\sim 10^{-2}$ AU.  It is worth mentioning that the power spectrum of irregularities in the direction of B1929+10 appears relatively steep (Bhat, Gupta \\& Rao 1999) and hence the refractive effects can be expected to be dominant, relatively speaking.  Also, the multipath scale (defined by the apparent size of the scatter-broadened image) when viewed in units of the Fresnel scale is an order of magnitude smaller for B1929+10 than that for B0950+08, indicating correspondingly weaker and less multiple scattering.  However, it is not clear as to how much of this apparent difference is due to the possible differences in the fractional distance to the scatterer (i.e. the scatterer location may be at one end in the case of B1929+10 instead of mid-way as we have assumed). It should be emphasised here that the refractive effects are highly episodic. So our ({\\it null}) results on these two pulsars, and the implied constraints (as in Figure \\ref{fig:ne_vs_a}), while being indicative, should be viewed as applying strictly only to the epochs of our observations.  \\begin{figure} \\epsfig{file=Cd131_f3.ps,width=6cm,angle=-90} \\caption[]{Upper bounds to the structure function of dispersion measure. The solid line is for PSR B0950+08, and the `dash' line is for B1929+10. See text for details.} \\label{fig:ne_vs_a} \\end{figure}   To summarise, radio signals emitted by very near-by sources may be scattered by discrete refractive irregularities that alter the diffraction patterns significantly, or may undergo `non-multiple' diffractive scattering.  Finding distinct spectral signatures of such events allows us to study the size and the density contrast of such large-scale irregularities. We have looked for such signatures in the signals from PSR B0950+08 and PSR B1929+10, but found no corresponding ACF feature significantly above the noise threshold ($7\\sigma$) of $2\\times 10^{-3}$ and $1.3\\times 10^{-3}$, respectively. The absence of any such signature/event in the observed data suggests useful limits on the structure function of dispersion measure corresponding to a wide range of refractive scales in the two lines of sight (as sampled at the epoch of our observations)."
    },
    "0011/astro-ph0011093_arXiv.txt": {
        "abstract": "We present new CO (1 $-$ 0) observations of eleven extragalactic tails and bridges in nine interacting galaxy systems, almost doubling the number of such features with sensitive CO measurements. Eight of these eleven features were undetected in CO to very low CO/HI limits, with the most extreme case being the NGC 7714/5 bridge.  This bridge contains luminous H~II regions and has a very high HI column density (1.6 $\\times$ 10$^{21}$ cm$^{-2}$ in the 55$''$ CO beam), yet was undetected in CO to rms T$_R$$^*$ = 2.4 mK.  The HI column density is higher than standard H$_2$ and CO self-shielding limits for solar-metallicity gas, suggesting that the gas in this bridge is metal-poor and has an enhanced N$_{H_2}$/I$_{CO}$ ratio compared to the Galactic value. Only one of the eleven features in our sample was unambiguously detected in CO, a luminous HI-rich star formation region near an optical tail in the compact group Stephan's Quintet. We detect CO at two widely separated velocities in this feature, at $\\sim$6000 km s$^{-1}$ and $\\sim$6700 km s$^{-1}$.  Both of these components have HI and H$\\alpha$ counterparts. These velocities correspond to those of galaxies in the group, suggesting that this gas is material that has been removed from two galaxies in the group. The CO/HI/H$\\alpha$ ratios for both components are similar to global values for spiral galaxies. ",
        "introduction": "Large amounts of interstellar gas can be removed from the main disks of spiral galaxies by four main processes: tides due to the gravitational force of a companion, ram pressure stripping during a near-head-on collision between gas-rich galaxies, ram pressure stripping by intracluster gas, and galactic winds driven by supernovae. By ejecting processed gas into intergalactic space, these mechanisms contribute to the metal-enrichment of the intergalactic medium. In many cases, which of these four processes is active in a given galaxy system can be determined from the optical, radio, and/or X-ray morphology. The signature of a tidal encounter between two galaxies is the presence of long stellar and/or HI tails and bridges (e.g., Toomre $\\&$ Toomre 1972), while head-on collisions between gas-rich galaxies can produce ring galaxies (e.g., Lynds $\\&$ Toomre 1976; Theys $\\&$ Spiegel 1977) as well as gaseous bridges between the galaxies (Struck 1997). Ram pressure stripping by intracluster gas leads to an HI deficiency (e.g., Giovanelli $\\&$ Haynes 1983), truncation of the outer HI disk of a galaxy (Cayette et al. 1990), and in some cases, bending of the HI disk (Kenney $\\&$ Koopmann 1999), but not usually removal of large quantities of molecular gas (Kenney $\\&$ Young 1989). Galactic winds are identified by extended extra-disk ionized gas without stellar counterparts (e.g., Rand, Kulkarni, $\\&$ Hester 1990).  Star formation sometimes occurs in gas clouds far removed from the main disks of galaxies. Of the four processes that can remove gas from disks, the best-known to trigger the formation of young stars in the stripped gas is tidal forces: numerous examples of luminous H~II regions in tidal features have been found (e.g., Schweizer 1978; Mirabel et al. 1991, 1992). Young stars have also been found in extra-disk gas clouds thought to have been stripped during head-on collisions (Smith et al. 1999) and interstellar-intracluster encounters (Xu et al. 1999).  In order to better understand the processes that lead to the removal of gas from galaxies and the triggering of star formation in this gas, it is important to make a complete inventory of the gas in these features.  This means not just the atomic gas, which has been surveyed in a large number of extragalactic tails and bridges (e.g., Haynes et al. 1984; Smith 1991; Hibbard $\\&$ van Gorkom 1996), but also the molecular gas. Molecular gas has proved elusive in classical tidal tails; in our earlier CO survey of six tidal tails (Smith $\\&$ Higdon 1994), no CO was found to very low levels, while only very low mass concentrations of molecular gas ($\\le$10$^7$ M$\\sun$) were found in the tidal features of the nearby interacting system M81/M82/NGC 3077 (Brouillet, Henkel, $\\&$ Baudry 1992; Walter $\\&$ Heithausen 1999).  In contrast to tidal features, gas removed from galaxy disks by ram pressure during head-on collisions between two gas-rich galaxies may be richer in CO. The proto-type of this class of object is the CO-rich gas concentration found outside the disk of the Virgo Cluster galaxy NGC 4438 by Combes et al. (1988). If the standard Galactic N$_{H_2}$/I$_{CO}$ ratio holds in this feature, it contains M$_{H_2}$ $\\sim$ 10$^9$ M$\\sun$. The ring-like morphology of NGC 4438 and the proximity of the companion galaxy NGC 4435 suggests that this clump was removed during a head-on collision between the two galaxies (Kenney et al. 1995). Another apparent example of gas ram pressure-stripped during a galaxy-galaxy collision is the eastern tail of the peculiar galaxy NGC 2782.  This feature has strong CO emission, corresponding to 6 $\\times$ 10$^8$ M$_{\\sun}$, assuming the Galactic N$_{H_2}$/I$_{CO}$ conversion factor (Smith et al. 1999). Based on morphological considerations, we surmised that this feature was created during a near head-on collision between two galaxies, rather than a grazing encounter (Smith 1994).  These results suggest that features produced in head-on collisions may differ in a fundamental way from classical tidal tails and bridges formed during grazing encounters: they may be richer in CO. During a head-on collision, the material pulled out into a tail or bridge may originate in the inner disk of one of the galaxies, and so may be more metal-rich than gas pulled out from the outer disk in a more distant encounter. The metallicity of the gas may affect the N$_{H_2}$/I$_{CO}$ ratio, and therefore the detectability of the CO line. Both theoretical (e.g., Maloney $\\&$ Black 1988) and observational (e.g., Wilson 1995; Verter $\\&$ Hodge 1995; Arimoto, Sofue, $\\&$ Tsujimoto 1996) studies show that low metallicities can lead to enhanced N$_{H_2}$/I$_{CO}$ ratios compared to the Galactic value. In gas with low abundances and dust content, ultraviolet radiation penetrates more deeply into a molecular cloud, causing a larger C$^+$ region relative to the CO core, increasing the N$_{H_2}$/I$_{CO}$ ratio.  The conclusion that tidal gas differs from ram pressure stripped gas is quite uncertain, in part because of the very small sample size. We have therefore continued our CO survey of extragalactic tails and bridges, selecting systems with high HI column densities, strong star formation rates, and/or ring-like morphologies. In this paper, we present new CO data for tails and bridges in nine additional interacting galaxy systems.  We compare these data with the previous results discussed above, as well as the new CO data discussed in Smith (2000), Braine et al. (2000), and Gao et al. (2000).  Throughout this paper, we assume H$_o$ = 75 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have obtained new CO (1 $-$ 0) observations of eleven extragalactic tails and bridges, and compared these measurements with previously-published data for another thirteen features. Of these 24 structures, four have inferred M$_{H_2}$/M$_{HI}$ upper limits (assuming the Galactic N$_{H_2}$/I$_{CO}$ conversion factor) less than global values for spiral galaxies, consistent with those found for irregular galaxies. Three of these four CO-poor features are in the NGC 7714/5 system; the most extreme case is the star forming bridge, which has a very high HI column density (1.6 $\\times$ 10$^{21}$ cm$^{-2}$ in the 55$''$ CO beam), yet was undetected in CO, giving an inferred M$_{H_2}$/M$_{HI}$ upper limit of $\\le$0.063. These features may have enhanced N$_{H_2}$/I$_{CO}$ ratios compared to the Galactic value due to low metallicities.  Of the 24 structures in our combined sample, eight have been detected in CO. Three of these eight (the Stephan's Quintet extra-disk source, the NGC 4038/9 tail, and the Arp 245 tail) have inferred efficiencies of star formation (L$_{H\\alpha}$/M$_{H_2}$) consistent with global values for spiral galaxies.  In contrast, the L$_{H\\alpha}$/M$_{H_2}$ ratios for the eastern tail of NGC 2782 and the NGC 4438 clump are lower than typical for disk galaxies, suggesting inhibited star formation. In none of the features were our CO limits low enough to conclusively show L$_{H\\alpha}$/M$_{H_2}$ ratios as high as those found in dwarf galaxies.  Of the eight detected features, two (NGC 2782 and NGC 4438) have optical morphologies indicating head-on collisions.  The Stephan's Quintet feature may also have been formed in this manner.  In these systems, we suggest that metal-rich gas with an approximately Galactic N$_{H_2}$/I$_{CO}$ ratio was pulled from the inner parts of the galaxies.  A few recent studies have reported the detection of CO in a number of classical tidal features (e.g., NGC 4038/9, Arp 245, and NGC 3561). In these structures, the HI column density may be locally high enough that the CO self-shielding limit is exceeded, in spite of modestly enhanced N$_{H_2}$/I$_{CO}$ ratios. There is a slight tendency for these tidal features to have lower M$_{H_2}$/M$_{HI}$ ratios and higher L$_{H\\alpha}$/M$_{H_2}$ ratios than ram pressure stripped features discussed above, but there is a large amount of scatter and the sample size is small. More observations are needed to determine whether there is truly a statistical difference in the gaseous and star formation properties of tidal and ram pressure stripped features.  \\vskip 0.1in  We are grateful for the help of the NRAO 12m telescope operators and staff in obtaining these data. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory under contract with NASA.  We are pleased to acknowledge partial funding for this project from a NASA grant administrated by the American Astronomical Society and from NSF grant INT-9908542."
    },
    "0011/hep-ph0011335_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is an amazing fact of our times that even as our understanding of cosmology progresses by leaps and bounds, we remain almost completely ignorant about the nature of most of the matter in the universe. According to recent fits to cosmological parameters \\cite{fit}, dark matter of some sort makes up close to 30\\% of the total energy density. This is much more than what is inferred from inventories of the luminous matter we can see. Moreover, the successes of big bang nucleosynthesis suggest that only a fraction of this matter can be made of ordinary baryons, like massive compact objects, faint stars, {\\it etc.}. Thus, unless gravity undergoes some drastic changes at distances larger than a few kpc (which is quite improbable from several points of view), we must postulate the existence of enormous amounts of dark matter of an unknown, non-baryonic origin.  A simple argument shows why current theories of particle physics are so prolific in suggestions for what the nature of this dark matter might be \\cite{GJK}. The vast majority of the proposed alternatives to the Standard Model involve new particles having masses which are of order $\\mw \\sim 100$ GeV, and which couple with electroweak strength to the ordinary matter which we know and love. If any of these particles is stable enough to have a lifetime as long as the age of the universe, it makes a natural candidate for dark matter. It does so because its abundance is naturally of the right order of magnitude so long as its interaction cross sections have weak-interaction strength. The abundance comes out right because it is set by the annihilation rate for particles which are initially in thermal equilibrium with ordinary matter. Cosmologically interesting abundances follow pretty much automatically for particles whose mass is of order $\\mw$ and whose annihilation cross sections have weak (or rather milliweak) interaction strength. (We give this argument in more detail within the body of the paper.)  Better yet, supersymmetric models, which are perhaps the best motivated of the many theories which have been proposed, very often have such long-lived states, due to the natural existence there of a conserved quantum number, $R$-parity, which keeps the lightest $R$-odd state from decaying. These particles cry out for interpretation as dark-matter particles, and it is no surprise that these models are by far the most widely explored in the literature \\cite{susyguys}.  Best of all, this explanation of the nature of dark matter can be tested experimentally. This is the direct goal of dedicated dark-matter detectors \\cite{Cad}, and an indirect goal of accelerator searches for events with missing energy, showing that a weakly-interacting particle has escaped the detector. If Nature smiles on us we soon may be treated to the discovery of new physics in both of these kinds of experiments. Indeed, recently the DAMA collaboration has announced the detection of a dark matter signal, as indicated by their seeing an annual modulation of the counting rate in a NaI detector \\cite{DAMAm}. However, the comparably precise data from the Ge detectors of the CDMS collaboration \\cite{CDMS} do not support these findings. (These two experiments need not be in contradiction with each other if the spin-dependent part of the cross-section is enhanced relative to the spin-independent part \\cite{PV}.)  Our goal in the present paper is to present a slightly unorthodox view. Although very well motivated, supersymmetric models are very complicated and enjoy an enormous parameter space. This makes them unable to definitively predict what dark-matter detectors must see. Furthermore, unlike the extensive evidence for the existence of dark matter, the arguments in favour of supersymmetry are almost exclusively theoretical. In our opinion, with the advent of good-quality data from dark-matter detectors, it behooves theorists to propose simple models for the dark matter which are consistent with present evidence, but which make definite predictions and so are easily falsifiable. These provide benchmarks against which other models and the data can be compared. We believe that it is only by comparing the implications of such models with one another, and with supersymmetry, that one can hope to properly interpret the data.  The model we study in this paper was first introduced by Veltman and Yndurain \\cite{Velt} in a different context. Its cosmology was later studied by Silveira and Zee \\cite{Zee}, and (with a complex scalar) by McDonald \\cite{McD}. It is the absolute minimal modification of the Standard Model which can explain the dark matter. It consists of the addition of a single spinless species of new particle, $S$, to those of the Standard Model, using only renormalizable interactions. To keep the new particle from interacting too strongly with ordinary matter, it is taken to be completely neutral under the Standard Model gauge group. Besides involving the fewest new states, the model is also just complicated enough to offer interestingly rich dark matter properties. Unlike the case if only spin-half or only spin-one singlet particles are added, it is possible for a singlet scalar to have {\\it both} significant renormalizable self-interactions {\\it and} renormalizable interactions with some Standard Model fields.  There is also a sense in which the model we propose is generic, should the dark matter consist of a single species of spinless particle. To this end, it is useful to ask the question of what a generic dark-matter model should look like. It is clear that the main property which one needs to ensure is the stability of the new particle, suggesting that the fields $S_{i}$, representing these particles appears in the Lagrangian in even powers, so that its decay is forbidden. If this field $S_{i}$ is considerably lighter than the rest of the other exotic undiscovered particles, these may be integrated out, leaving an effective Lagrangian at electroweak scale which has the generic form \\begin{equation} {\\cal L}={\\rm kinetic~term~for}~S_{i}+S_{i}^{*}MS_{i}+\\sum S_{i}^{*}O_{\\sss SM}S_{i}+... \\end{equation} where the kinetic, mass terms and interactions with the SM (via the set of operators $O_{\\sss SM}$) in general would depend on the spin of $S_{i}$. The most important couplings at low energies are those of lowest dimension, corresponding to the lowest-dimension choices for the operators $O_{\\sss SM}$. Our model also has this form, with only a single singlet scalar $S$. In this language, our dropping of all nonrenormalizable interactions corresponds to keeping only those interactions which are consistent with (but do not require) all other exotic particles to be arbitrarily heavy compared with the weak scale. We might expect our model to therefore capture the physics of any more complicated theory whose impact on the dark matter problem is conveyed purely through the low-energy interactions of a single spinless particle.  An additional, more tentative, incentive for formulating more models stems from recent indications of problems with subgalactic structure formation within the non-interacting cold-dark-matter scenario \\cite{problem}. A `generalized' form of cold dark matter may avoid these problems if its self-interactions\\footnote{ The required self-interactions however also lead to spherical halo centers in clusters, which are inconsistent with the ellipsoidal centers indicated by strong gravitational lensing data \\cite{Yoshida}.} can produce scattering cross sections of order $10^{-24}$ cm${}^2$ in size \\cite{selfint}. Within the present context this proposal would require the masses of dark matter particles not to exceed 1 GeV. Within the minimal model described in this paper, we find this range of masses may be just barely possible, but requires unnatural fine tunings due to the relationship between masses and couplings imposed by the maintenance of the correct cosmic abundance of dark-matter scalars.  This paper is organized as follows. In the next section we identify the three parameters which describe the model, and determine the general conditions which lead to acceptable masses and to sufficiently stable dark matter. In section 3 we calculate the annihilation cross section of $S$-particles and give the resulting cosmic abundance as a function of masses and couplings. This calculation is similar to the analysis of Ref. \\cite{McD}. We perform the numerical analysis for the most interesting part of the parameter space, with 100 GeV $\\le \\mh\\le$ 200 GeV and 10 GeV $\\le \\ms \\le $ 100 GeV. In section 4 we obtain the  cross section for elastic scattering with ordinary matter and apply the constraints, imposed by direct and indirect searches. Section 5 computes the cross sections for the missing energy events which are predicted for colliders due to the pair production of $S$ particles. It also contains a prediction for the degradation of the Higgs boson signal at hadronic colliders, when the Higgs boson is allowed to decay into a pair of $S$ particles. Our conclusions are reserved for section 6. ",
        "conclusions": "We have presented the first study of the minimal model for non-baryonic dark matter, which consists of the Standard model plus a singlet scalar. This model is characterized by one additional real scalar field and three new parameters. The absence of linear and cubic terms in $S$ is required to ensure that the the new scalar is sufficiently stable to contribute significantly to the dark matter currently present in the Universe. This stability also precludes the development of a nonzero v.e.v. for $S$. The study of the model's scalar potential shows that electroweak symmetry is spontaneously broken while $\\langle S \\rangle$ vanishes for a significant domain of parameter space.  At the renormalizable level the three parameters describe the $S$ particle mass, its self-coupling and its couplings to other Standard Model fields, which all are mediated by a coupling to the Standard Model Higgs, of the form $\\l S^{2}(H^{\\dagger }H)$. The simplicity of this coupling significantly simplifies the calculations of primordial abundances and observable signals at colliders and underground detectors, making the resulting predictions more certain. The primordial abundance is governed mainly by the annihilation of $S$-particles via $s$-channel Higgs exchange. If the mass of the Higgs is not too close to $2\\ms$, the observed abundance of dark matter is achieved by a most natural choice for the coupling, $\\lambda \\sim O(0.1 - 1)$.  These large values of $\\lambda$ lead to nuclear elastic cross sections which are $ \\sim 10^{-43}$ cm$^2$  in size. This is slightly below the limits of sensitivity of the DAMA and CDMS experiments, but is likely to be detectable at future experimental facilities. Thus, the direct detection of $S$ particles is not yet ruled out, and is easily feasible within the near future! Our minimal model is unable to reconcile the DAMA and CDMS results, and so predicts that one or the other is incorrect.  Cross sections this large also lead to a significant flux of high-energy neutrinos generated by $S$ annihilation at the centers of the earth and the sun, at levels potentially accessible to large neutrino detectors. We give here only an estimate of the expected flux, and we believe the encouraging results motivate further work to simulate the expected neutrino and muon spectra and intensities. It would also be intriguing to repeat the numerical analysis of ref. \\cite{McD}, updated with the Higgs masses favored by modern collider results.   Large values for $\\l$ may also lead to the significant missing-energy signals at colliders, corresponding to the production of a pair of $S$-particles. This signal is unlikely to be seen at hadronic machines, where the background events (two jets plus missing energy) will have much larger cross-sections. Lepton colliders, such as LEP or NLC are needed. The possible existence of $S$-particles with a mass smaller than half of the Higgs boson mass poses a significant threat to Higgs searches at the Tevatron and LHC. Indeed, in this case the invisible decay of the Higgs boson can be many orders of magnitude larger than the search modes, making the Higgs effectively invisible. For intermediate Higgs masses, up to 140 GeV and higher, the pair production of $W$ bosons introduces a nonnegligible visible width and so saves the Higgs searches. It is important that in the dangerous domain of parameter space, $2 \\ms < \\mh < 140$ GeV, the size of the elastic cross-section, $\\sigma_{\\rm el}$, is limited {\\em from below} to be larger than $10^{-44}$ cm$^2$. If this were indeed the case, the next generation of the underground detectors would likely discover the recoil signal.  Because the self-coupling, $\\ls$, is unconstrained by the abundance condition, one might hope to use models of this sort to produce a strongly-interacting dark matter candidate, a possibility recently advocated in ref.~\\cite{selfint,BBRT}. Unfortunately, this proposal would also require a rather small mass for the $S$-scalars (1 GeV or less), and since $\\l$ cannot be small for such masses, one is led to a fine-tuned (ppm) cancellation between $\\lambda v_{EW}^{2}$ and the ``bare'' mass $m_{0}^{2}$. Worse, having $\\ms$ be 500 MeV or smaller pushes $\\l$ to the very limits of our perturbative analysis.  It remains an interesting challenge to construct a non-minimal model of strongly-interacting dark matter along the present lines. One possibility is to suppress $\\l$ to evade the fine-tuning problem, and set the freeze-out abundance by a set of non-renormalizable operators of dimension 6 and higher.  Another possibility for having light and strongly interacting scalar dark matter requires $\\l$ to be very small, scalar particles to be out of equilibrium at all temperatures, and their abundance to be fine-tuned (for example,  using inflationary scenarios). The suppression of the coupling to the Higgs, strong self-interaction and a low mass for $S$ could be simultaneously achieved in models where $S$ is composite \\cite{FP}, in which case there is unlikely to be significant impact on Higgs physics.  Our pursual here of a more model-independent approach to the physics of dark matter particles is ultimately stimulated by the maturity of current dark-matter detection efforts. As the data erodes the large parameter space of the popular neutralino models, or if a dark matter candidate is eventually found, a reliable interpretation of the experiments can only be found by comparing with a wide variety of explanatory models. Our proposal in this paper marks the absolute minimum of complexity required by such a model, and from this minimality springs its predictiveness. It is also generic in the sense that it is the most general renormalizable model consistent with the assumed particle content and symmetries. As such, it captures the implications of any models whose implications for dark matter arise from a low energy spectrum which contains only these particles and symmetries. We believe that further model--building efforts are in order as experimental developments progress.  \\bigskip  \\bigskip  \\bigskip  \\centerline{\\bf Acknowledgements} \\bigskip We thank G. Couture, T. Falk, L. Kofman, G.Mahlon, K. Olive, S. Rudaz and M. Voloshin for usefull discussions and comments. M.P. wishes to thank the Physics Departments of McGill University and the University of Qu\\'ebec at Montr\\'eal, where a significant portion of this work was carried out. M.P. and T. t V.'s research was funded by the US Department of Energy under grant number DE-FG-02-94-ER-40823, and C.B. acknowledges the support of NSERC (Canada), FCAR (Qu\\'ebec) and the Ambrose Monell Foundation.  \\newpage   \\def\\NPB#1#2#3{{\\it Nucl. Phys.} {\\bf B{#1}} (#2) #3} \\def\\PLB#1#2#3{{\\it Phys. Lett.} {\\bf B{#1}} (#2) #3} \\def\\PRD#1#2#3{{\\it Phys. Rev.} {\\bf D{#1}} (#2) #3} \\def\\PRL#1#2#3{{\\it Phys. Rev. Lett.} {\\bf {#1}} (#2) #3} \\def\\PRT#1#2#3{{\\it Phys. Rep.} {\\bf {#1} C} (#2) #3} \\def\\MODA#1#2#3{{\\it Mod. Phys. Lett.} {\\bf {#1}} (#2) #3}"
    },
    "0011/hep-ph0011279_arXiv.txt": {
        "abstract": "{ Low energy string theory predicts the existence of an axion field which can lead to cosmic birefringence. We solve the electromagnetic wave equations in the presence of such an axion and a dilaton field in order to determine their effect on the polarization of light. We find that the presence of dilaton field leads to a nontrivial modification of the final result. We comment on the possibility of discovering such an effect by observations of radio wave polarizations from distant radio galaxies and quasars. We have also determined the limits on the string theory parameters that are imposed by the current radio polarization data. } ",
        "introduction": "There have been several studies which explore the possible existence of cosmic birefringence \\cite{bire}. In recent work Kar et al \\cite{Kar} and Majumdar and Sengupta \\cite{Partha} has pointed out that the Kalb-Ramond field, which arises in supergravity theories, makes the space time birefringent. The authors solve the electromagnetic wave equations in the presence of such a field and predict that this field will lead to a redshift dependent rotation of polarization whose details depend on whether the universe is radiation or matter dominated. Earlier in Ref. \\cite{CFJ} Carroll et al analyzed the polarizations of radio waves from distance galaxies and quasars as a function of redshifts and found no statistically significant effect. The observable of interest in that study was the angle $\\beta=\\chi-\\psi$, where $\\psi$ is the orientation angle of the axis of the radio galaxy and $\\chi$ the observed polarization angle after the effect of Faraday rotation is taken out of the data. This can be done by making a straight line fit to the polarization angle $\\theta(\\lambda^2)$ which can be expressed as \\begin{equation} \\theta(\\lambda^2) = RM\\ \\lambda^2 + \\chi \\end{equation} where $\\lambda$ is the wavelength, $RM$ and $\\chi$ are the slope and the intercept of the fit. $RM$ stands for rotation measure and depends on the plasma density and the parallel component of the background magnetic field along the direction of propagation of the wave. Later Ref. \\cite{NR} claimed the existence of a large scale anisotropy in the same data. This claim was, however, questioned by many authors \\cite{EB,CF97,LFW}.  The possibility of local anisotropy, independent of redshift, was explored in Ref. \\cite{jain98} and earlier in \\cite{Birch,KY}. These authors found a statistically significant effect. The original claim made by Birch \\cite{Birch} in 1982 was dismissed by Bietenholz and Kronberg \\cite{BK} in 1984 who compiled a larger data set which they claimed did not show any effect. However Jain and Ralston \\cite{jain98} argued that Ref. \\cite{BK} did not pay careful attention to the parity of the correlation ansatz \\cite{RJ}. When the transformation properties of statistical distribution are taken into account the data does show a significant effect which is further enhanced by choosing an appropriate statistic as well as by putting some cuts on the data.  Inspite of the progress in our understanding in string theory in the last two decades, we are yet to have a string theory prediction that can be tested in the experiments. One of the main reason for such an unsatisfactory situation is that the energy scale at which stringy effects become significant is beyond the reach of present day experiments. It is thus believed that the most likely area for a confrontation between string theory and experiments is through extracting astrophysical or cosmological consequences of string theory which might be measurable today. With this kind of motivation in mind, we study in this paper the prediction of string theory for cosmic birefringence. However, as our knowledge of string theory is mostly perturbative in nature, we face several problems in carrying out such a study. These are mainly due to the presence of various massless moduli that are present in the low energy limit of string theory. We will point those out as we proceed.  Four dimensional heterotic string theory, in its low energy limit, contains  moduli coming from compactification, gauge fields, dilaton and an axion. In this theory, coupling among various fields are unique and they are all controlled by the expectation value of dilaton. Given such a framework, we analyse the propagation of electromagnetic waves in curved metric background. Due to the presence of dilaton, as we will argue, medium turns into a conducting medium where the time variation of dilaton plays the role of effective conductivity. On the other hand, the axion  rotates the plane of polarization of the electromagnetic wave causing non-trivial birefringence. However, notice that the dilaton (and other moduli) is a massless scalar in string theory. An important issue to resolve in string cosmology is to explain the absence of massless dilaton at the present time. One expects that a potential for dilaton will be generated in string theory through some non-perturbative effects and the dilaton will sit at the minimum of the potential picking up a mass. The mechanism of how precisely this should happen is still unknown in string theory. This, in turn, creates the most severe problem in comparing observational data with  string theory. In the rest of the paper, we continue to treat dilaton as a massless scalar and find string prediction for cosmic birefringence. We hope that some of our results will be helpful when dilaton decoupling in string theory is properly understood. In general, we find that the relevant equations are hard to solve exactly. However, it is well-known that heterotic string theory in four dimensions is self dual. The duality group is parametrized by $SL(2,R)$ matrices. Using this property, in the next section, we  generate some non-trivial solutions for various fields starting from some simple seed configuration. This, in turn, allows us to calculate birefringence of electromagnetic waves, propagating in radiation or matter dominated universe, in a fairly straightforward manner. However, since the solution generating technique does not saturate all possible allowed configurations, in section III, we solve the electromagnetic equations within the WKB approximation. We then discuss our results in section IV. ",
        "conclusions": "In this paper, we have studied cosmic birefringence in the framework of perturbative four-dimensional heterotic string theory. Due to coupling between dilaton and axion in this theory, we have shown that the dilaton field changes the birefringence produced by the axion field in a nontrivial way. We obtained an approximate expression for the rotation in polarization predicted by heterotic string theory within the WKB approximation. We analysed the data for distant quasars and radio galaxies to determine if the effect is present in the data. We find a marginally significant effect. However the correlation is caused primarily by two largest redshift points and the effect is lost after these points are eliminated. Hence we find that the data does not show significant correlation but the large redshift points are suggestive that a signal might exist. This should be thoroughly investigated by accummulating more data at the largest redshifts. We have also determined the limits on the string theory parameters that are imposed by the current polarization data from distant quasars and radio galaxies.  We, furthermore, like to point out that all these calculations were done taking dilaton as a massless scalar. As explained in the introduction, one of the biggest problem in string theory is to find a mechanism to remove this massless dilaton from present day physics. One expects that a potential for dilaton will be generated in string theory through some non-perturbative effects and the dilaton will sit at the minimum of the potential picking up a mass. As the mecanism of how precisely this should happen is still not clear in string theory, we have not dealt with such a scenario in this paper.  \\bigskip \\noindent {\\bf Acknowledgements:} PJ thanks John Ralston for very useful discussions and for providing very useful insights into the statistical analysis of this data. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
    },
    "0011/astro-ph0011266_arXiv.txt": {
        "abstract": "We present a photometric method for identifying stars, galaxies and quasars in multi-color surveys and estimating multi-color redshifts for the extragalactic objects. We use a library of $> 65000$ color templates for comparison with observed objects. The method was originally developed for the Calar Alto Deep Imaging Survey (CADIS), but is now used in a variety of survey projects. We checked its performance by spectroscopy of CADIS objects, where it provides high reliability (6 mistakes among 151 objects with $R<24$), especially for the quasar selection, and redshifts accurate within $\\sigma_z \\approx 0.03$ for galaxies and $\\sigma_z \\approx 0.1$ for quasars. For an optimization of future surveys, a few model surveys are compared, which use the same amount of telescope time but different filter sets. In summary, medium-band surveys perform superior to broad-band surveys although they collect less photons. A full account of this work is already in print ([11,12]). ",
        "introduction": "The value of a multi-color survey is usually not simply limited by the depth of its imaging but by the depth of its classification used to select the objects of interest from the survey data. Here, we investigate what factors the depth of a successful classification depends on. Obviously, it depends in some way on the filter set used as well as on the depth of the imaging, but the optimum choices might vary according to the goal of the survey.  If a survey aims at identifying only one type of object with characteristic colors, a tailored filter set can be designed. E. g., when looking for U-band dropouts [9], the {\\it ugr} filter set is a very good choice. The performance of such a dropout survey depends mostly on the depth reached in the U-band, so the photon flux detection limit in U is the key figure. Also, number count studies are limited by the completeness limit in the filter of concern.  If all kinds of different objects are the subject of the survey, a filter set covering the entire spectral range accessible to the survey instrument is a natural choice. The SLOAN Digital Sky Survey [13] is presently the most ambitious project to provide a broad-band {\\it ugriz} color database, on which the astronomical community might perform a large number of virtual surveys.  Of course, it is possible to split the bands of a filter set even further into a larger number of medium-band filters. Surveys like CADIS with typically 10 to 20 filters are sampling the visual spectrum with a resolution almost comparable to that of low resolution imaging spectroscopy. Intuitively, one might expect that such surveys are very expensive in terms of telescope time while their particular value might not be obvious. This paper will show, that in fact medium-band surveys are a powerful and economically competitive alternative to broad-band surveys. ",
        "conclusions": "We presented an innovative method that performs a multi-color classification and redshift estimation of astronomical objects in a unified approach. The method is essentially based on templates and evaluates the statistical consistency of a given measurement with a database of spectral knowledge, serving as a second, very crucial input to the algorithm. Spectroscopic confirmation within CADIS showed that the method works well and is of high practical relevance.  A simulation of a few model surveys with fundamentally different filter sets revealed that in practice surveys with many medium-band filters are superior to pure broad-band surveys for the purpose of classification and redshift estimation. Even when being limited by telescope time, the medium-band filter sets achieve at least the same working depth as the broad-band survey while providing much improved redshift resolution.  We conclude, that this method should be very suitable for many survey applications, which usually require only low spectral resolution and finite accuracy in the derivation of physical parameters, but aim for large samples to feed statistical studies and to search for rare and unusual objects. Of course, if a 100\\% sure confirmation on the nature of an individual object is needed, or if you aim for high resolution studies, the method gives only a preselection of candidates."
    },
    "0011/astro-ph0011050_arXiv.txt": {
        "abstract": "The enrichment of Pb in the Galaxy is followed in the framework of a detailed model of Galactic chemical evolution that already proved adequate to reproduce the chemical enrichment of O and of the elements from Ba to Eu. The stellar yields are computed through nucleosynthesis calculations in the Asymptotic Giant Branch (AGB) phase of low- and intermediate-mass stars, covering a wide range of metallicities. The physical parameters of the stellar structure were derived from full stellar evolutionary models previously computed. We show that low-mass AGB stars are the main producers of Pb in the Galaxy, with a complex dependence on  metallicity and a maximum efficiency  at [Fe/H] $\\sim$ $-$1. Our calculations succeed in reproducing the abundances of Pb isotopes in the solar system: the role attributed by the classical analysis of the $s$ process to the {\\it strong component}, in order to explain more than 50\\% of solar $^{208}$Pb, is actually played by the high production of Pb in low-mass and low-metallicity AGB stars. We then follow the Galactic chemical evolution of Pb isotopes and give our expectations on the $s$-process contribution to each of them at the epoch of the solar system formation. Finally, we present new spectroscopic estimates of Pb abundance on a sample of field stars and compare them, together with a few other determinations available, with the predicted trend of [Pb/Fe] in the Galaxy. ",
        "introduction": "Neutron captures are the main way for the building up of the nuclei heavier than iron. According to the seminal work by Burbidge et al.~(1957), in order to reconstruct the Galactic evolutionary history of heavy elements one has to consider two major mechanisms of neutron addition: the $s$-process (slow neutron-captures) and the $r$-process (rapid neutron-captures). For the $s$-process, the abundance distribution in the solar system was early recognized to derive from a non unique site. Indeed, the phenomenological approach of $s$-process studies, or {\\it classical} analysis, was shown to require three components to account for the abundances of all $s$-nuclei from Fe to Bi: the {\\it weak}, the {\\it main}, and the {\\it strong} component (Seeger, Fowler, \\& Clayton~1965; Clayton \\& Rassbach~1967; K\\\"appeler et al.~1982; K\\\"appeler, Beer, \\& Wisshak~1989).  The weak component, responsible for the $s$-process nuclides up to $A \\simeq$ 90, was ascribed to neutron captures occurring in advanced evolutionary phases of massive stars, through the activation of the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction (Lamb et al.~1977; Prantzos et al.~1990; Raiteri et al.~1993 and references therein). The main component, feeding the heavier $s$-process nuclides up to $A =$ 208, is known to originate in AGB stars of Galactic disk metallicity during recurrent thermal pulses, hereafter TP. Many theoretical and observational works (see Busso, Gallino, \\& Wasserburg~1999 for a review) converge on the idea that Low-Mass Stars (hereafter LMS, $M$ $\\le$ 4 \\ms) play the dominant role in $s$-processing, and that  neutrons are released in radiative conditions in the interpulse phase via the $^{13}$C($\\alpha$,n)$^{16}$O reaction (Straniero et al.~1995). A tiny $^{13}$C {\\it pocket} is thought to develop in the top layers of the He- and C-rich zone (called He intershell) as a consequence of the penetration of a small amount of protons from the envelope, to burn locally and to produce neutrons before the  development of the next TP. During advanced TPs a second small neutron burst is released by the marginal activation of the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction. Finally, the strong component was advanced on theoretical grounds by the classical analysis of the $s$-process (Clayton \\& Rassbach 1967), in order to reproduce more than 50\\% of $^{208}$Pb, the most abundant Pb isotope in the solar system, although the astrophysical site for this process remained unknown.  In this work we study the Galactic enrichment of Pb in the light of AGB nucleosynthesis results presented by Straniero et al.~(1997) and Gallino et al.~(1998). One of the main goals of these works was the development of a new set of model calculations for AGB stars with different masses and metallicities, to provide stronger evolutionary supports to nucleosynthesis calculations. The results demonstrated in detail the dependence of the $s$-process yields on stellar metallicity, and introduced the idea that a separate astrophysical site for the strong component might be unnecessary (Gallino et al.~1998), because a large production of $^{208}$Pb derives from AGB stars of low metallicity. This expectation must however be substantiated with a  detailed model for the Galactic enrichment, to verify whether the production of Pb isotopes is really adequate.  The paper is organized as follows: in \\S~2 we discuss the sensitivity of Pb production (in particular $^{208}$Pb) to metallicity and to different assumptions on the $^{13}$C pocket in low-mass AGB stars. We also analyze the role of intermediate-mass AGB stars in the Pb production at different metallicities. In \\S~3 we present spectroscopic estimates by synthetic spectra of the Pb abundance on a limited sample of field metal-poor stars using the Pb line at 3683.48~\\AA~ and outline the difficulties and the aims of such an analysis. In \\S~4 we briefly introduce the Galactic Chemical Evolution (hereafter GCE) model adopted, present our predictions for the $s$-components of the Pb isotopes in the solar system, and compute the [Pb/Fe]\\footnote{We adopt the usual spectroscopic notation [Pb/Fe] = log (Pb/Fe) $-$ log (Pb/Fe)$_{\\odot}$} trend vs. metallicity in the three Galactic zones in which the GCE model is organized (halo, thick disk and thin disk). Then we compare our model predictions with the above new spectroscopic determinations and with a few other published data, outlining the need of a dedicate observational effort. Finally, in \\S~5 we summarize the main conclusions and point out a few aspects deserving further analysis. ",
        "conclusions": "In this paper we have studied the evolution of Pb in the interstellar gas of the Galaxy, adopting a new set of models for $s$-processing in AGB stars of different metallicities, and new observational constraints of unevolved field stars for Pb abundances.  We have calculated the $s$-process yields with post-process calculations based on AGB models and applied these yields in the framework of a Galactic model that follows the GCE of the halo, thick and thin disk. Taking into account the important role played by different generations of low-metallicity AGB stars for the production of Pb in the Galaxy, we obtain at $t = t_\\odot$ a Pb $s$-fraction  of 91\\%, together with a sharp increase at [Fe/H] $\\simeq -$1.5. Confirming the prediction by Gallino et al.~(1998), we also found that the production of the most abundant Pb isotope, $^{208}$Pb, is matched when the $s$-process occurring in low-metallicity AGB stars is properly considered in the context of the chemical evolution of the Galaxy. Only a small portion of $^{208}$Pb can be attributed to the $r$-process, even accounting for decays of transuranic isotopes.  {"
    },
    "0011/astro-ph0011099_arXiv.txt": {
        "abstract": "Using the {\\it Rossi X-ray Timing Explorer} and the Nordic Optical Telescope we have obtained the highest ever quality X-ray/white-light high-speed photometry of XB\\thinspace1916--053.  We refine the X-ray period ($P_X$) to $3000.6\\pm0.2$\\sec\\ via a restricted cycle counting approach.  Using our complete optical lightcurve, we have extended the optical period ($P_{opt}$) ephemeris by another 4 years, providing further evidence for its stability, although a slightly longer period of $3027.555\\pm0.002$\\sec\\ now provides a marginally better fit.  Moreover, modulations at both $P_X$ and $P_{opt}$ are present in the optical data, with the former dominating the nightly lightcurves (i.e. a few cycles of data).  We have also attempted to determine the ``beat'' period, as seen in the repeating evolution of the X-ray dip structure, and the variation in primary dip phase.  We find that a quasi-period of $4.74\\pm0.05$\\dy\\ provides the best fit to the data, even then requiring phase shifts between cycles, with the expected 3.90\\dy\\ ``beat'' of $P_X$ and $P_{opt}$ appearing to be less likely. Finally, considering the nature of each of these temporal phenomena, we outline possible models, which could explain all of the observed behaviour of this enigmatic source, focusing on which of $P_X$ or $P_{opt}$ is the binary period. ",
        "introduction": "The low-mass X-ray binary (LMXB) \\targ\\ was the first of the 10 dipping sources to be discovered.  The recurrence period of the dip \\til 3000~s \\cite{walt82,whit82a} also provided one of the first binary period determinations for a LMXB, and places it as the shortest period dipper.  It also exhibits type-I X-ray bursts confirming that the compact object is a neutron star.  A low-mass (\\til0.1\\msun)  donor is then required to fit into the system, which is most likely evolved, partially degenerate and hydrogen depleted \\cite{nels86}.  The lack of eclipses by this companion, for a representative radius \\til0.1\\rsun, yields an inclination in the range 60\\degree\\ to 79\\degree\\ \\cite{smal88}.  The generally accepted model for the origin of X-ray dips, suggests that they are due to the obscuration of the central source by vertically extended material at or close to the outer edge of the accretion disc \\cite{parm88,armi98}.  In recent years Church and co-workers have studied a selection of dippers, utilising the superior spectral capabilities of {\\it ASCA} and {\\it BeppoSAX}, to probe in detail the spectral evolution through the dips \\cite{chur93,chur95}. They find that a partial covering absorption model reproduces the observed spectra well, where the fractional contribution from each of two components (point-like blackbody and more spatially extended cut-off power law) changes independently.  For \\targ\\ a physical model is implied consisting of the point-like blackbody emission region (presumably the neutron star surface and/or adjacent boundary layer), plus the cut-off power law component, originating in an extended accretion disc corona (ADC, responsible for Comptonization, \\pcite{chur97}).   In the non-dip spectra the blackbody alone contributes  \\til30\\% to the total flux \\cite{chur97,chur98}.  During dips the blackbody component is very rapidly and totally obscured, whilst the power-law is more gradually attenuated, as the partial-covering fraction increases (even to complete coverage in the deepest dips seen). From measurement of the ingress/egress times a size for the radial extent of the ADC is found of 20\\% of the tidal disc radius.  \\begin{figure}[!htb] \\centering \\includegraphics[scale=0.41,angle=0]{Px_Popt_history.ps} \\caption{Determinations of the ``orbital'' period of \\targ.  The X-ray period ($\\sim$3000s) measurements from a long-line of satellite missions at the top and the three previous optical results at the bottom {(\\small 1. White \\& Holt 1982, 2. Walter {\\em et al.} 1982, 3. Smale {\\em et al.} 1989, 4. Yoshida {\\em et al.} 1995, 5. Church {\\em et al.} 1997, 6. Church {\\em et al.} 1998, 7. Morley {\\em et al.} 1999, 8. this paper, 9. Callanan, Grindlay \\& Cool 1995, 10. Grindlay {\\em et al.} 1988, 11. Grindlay 1989, 12. Schmidtke 1988).} \\label{fig:Pdets}} \\end{figure}  \\subsection{Multi-periodicities of \\targ} Following their discovery of an optical counterpart to \\targ\\ \\scite{grin88} searched for and found an optical modulation close to the 3000\\sec\\ dip recurrence period.  Subsequently, this source has been the object of many timing studies.  It was soon apparent that the periods derived from X-ray and optical observations were inconsistent, differing by \\til 1\\%.  The results of these determinations are summarised in Fig.\\ref{fig:Pdets}.  For the X-ray dips, the earliest method adopted (by \\pcite{walt82,whit82a}), was to consider the timings of a series of (near-consecutive) dips from a single satellite pointing, and calculate the best fit recurrence time.  Unfortunately, the short observing baseline leads to large uncertainties in the period. More recently, the approach generally taken has been to epoch-fold the lightcurves on a grid of trial periods, and compute chi-squares to identify the period yielding the most distinct modulation. However, given the notable variation in the morphology of the dip structure over time (see e.g. \\pcite{smal88,smal92b,yosh95,chur97}), the precision of this method is also uncertain, and the large errors quoted may still be underestimates.  In their optical studies, \\scite{schm88} and \\scite{grin88} used a straightforward Fourier transform of their long (and irregularly) sampled lightcurves, whilst \\scite{call95} refined this by way of cycle-counting using the minima in the folded optical lightcurves to define the ephemerides.  Not only are two close periodicities present, there is evidence for two longer periodic phenomena.  In fact, it is assumed that physically one of the two $\\sim3000$s periods arises from binary motion, whilst the other is due to a beat with another longer period effect, which must be close to 4 days.  Indeed changes in the lightcurve shapes in both bands and in particular the phases of dip features has been seen on this timescale \\cite{grin89,smal89,yosh95}.  Note, for brevity, we shall henceforth refer to any \\til4\\dy\\ period as the ``beat'' period of $P_X$ and $P_{opt}$, whilst reserving the notation $P_{beat}$ for the quantity calculated as $(P_X^{-1} -P_{opt}^{-1})^{-1}$, and using $P'_{beat}$ for any other period close to this, i.e. not assuming that they are the same.  Lastly, an analysis of {\\it Vela 5B} data by \\scite{prie83} yielded a possible longer-term variation in X-ray flux of 199 days.  But with a false alarm probability of 10-20\\% \\cite{smal92}, this detection was by no means firm.   A number of models have been proposed to explain the multi-periodicities, as we shall discuss in section~\\ref{sect:disc}, but no clear consensus has been achieved given the limited constraints imposed by the observations to date.  It was clear that further complementary X-ray and optical observations were vital in order to solve the mystery of the multiple periodicities and constrain the range of proposed models to explain them.  Hence, we obtained time with both the {\\it Rossi X-ray Timing Explorer (RXTE)} and the Nordic Optical Telescope (NOT) to undertake a simultaneous X-ray/optical campaign.  It is the results of these observations that we present here.  In section~\\ref{sect:ObsDR} we summarise the observations and data reduction performed, moving on to the temporal analysis in \\S~\\ref{sect:Tanal1} (X-ray period), \\S~\\ref{sect:Tanal2} (``beat'' period), \\S~\\ref{sect:Tanal3} (longer-term variations) and \\S~\\ref{sect:Tanal4} (optical periods).  We discuss our results in section~\\ref{sect:disc}. ",
        "conclusions": "} \\begin{tabular}{l l c l } % Temporal phenomena & Analysis undertaken & Result & Notes\\\\ % Recurring X-ray dips &PDM of entire X-ray (PCA) dataset & $P_X=3001.0\\pm4.0\\sec$ & Drops in flux from non-dip level of varying \\\\ &\tCycle counting between 6 narrow & & duration, but all within narrow phase range.\\\\ &\t dips (over 3 month span) &$P_X=3000.6\\pm0.2\\sec$  &\\\\ Smooth modulation & Cycle counting between minima& \t$P_{opt}=3027.555\\pm0.002\\sec$&  Smooth, almost sinusoidal variation\\\\ of optical flux\t\t\t\t\t\t& over 1987--1998&&but on short timescales\\\\ &&&dominant features not phase stable\\\\ &&& (follow X-ray period  instead)\\\\ ``Beat'' period & $P_{beat}=(P_X^{-1}-P_{opt}^{-1})^{-1}$  & $P_{beat}=3.90\\pm0.03\\dy$ & - \\\\ Dip morphology cycle & Sinusoid fitting to dip phase variation + &  $P_{beat}'=4.74\\pm0.04\\dy$& Dip phase variations approximated by sinusoid,\\\\ &\tconsistency with systematic/cyclic evolution&&but evidence for quasi-periodic nature\\\\ Longer-term flux& LSP/PDM of 1 day average &\t $P_{long}=83.0\\pm0.4\\dy$ & Smooth variation (though non-sinusoidal)\\\\ modulation \t\t\t\t\t\t\t\t\t\t\t\t&\tX-ray (ASM) flux levels\t&& Energy independent amplitude\\\\  \\end{tabular}  \\end{sidewaystable} \\newpage"
    },
    "0011/astro-ph0011116_arXiv.txt": {
        "abstract": "Re-analysis of some archival X--ray data of black hole binaries in soft spectral states is presented. Results of detailed spectral fitting, including proper physical description of the X--ray reprocessed component, indicate that in most cases the soft, thermal disk emission cannot be described by a simple multi-color blackbody model. Additional Comptonization of those seed photons provides an acceptable description of the soft component, implying a presence of a warm ($k\\Te \\sim 1$ keV), mildly optically  thick ($\\tau \\sim 10$) phase of accreting plasma.  Combined spectral/temporal analysis of GS~1124-68 during 'flip-flop' transitions in VHS suggests that the optically thick accretion flow can proceed in (at least) two modes giving the same time-integrated behaviour (i.e.\\ energy spectra) but very different timing behaviour. ",
        "introduction": "Accreting black holes, both stellar (including the Galactic $\\mu$--quasars) and supermassive, show strong, soft, thermal components in their EUV--X-ray spectra, when accreting above certain fraction of Eddington rate. This component is thought to come from an optically thick accretion disk, and its spectrum is usually modeled as a simple, multi-color blackbody. For re-examining a number of such spectra, I will use the spectral models described in \\citeauthor{ZDS98} \\shortcite{ZDS98}.  \\subsection{The Intermediate State}  An example of IS is the May 18th observation of GS~1124-68 (Nova Muscae 1991). Using first the frequently employed but unphysical model, {\\sc diskbb + smedge(powerlaw)}, gives an unacceptable fit, $\\chi^2_{\\nu} = 38/25$. Adding a Gaussian line at 6.4 keV does {\\em not\\/} improve the fit, unless the line width is $\\sim 10$ keV ($\\chi^2_{\\nu} = 28/23$). Allowing the line energy to be free gives $\\chi^2_{\\nu} = 21/23$ for $E = 4\\pm 0.5\\,$keV. These values have little to do with Fe spectral features, but rather indicate that it is the continuum shape that is not properly modeled. Using now the {\\sc diskbb + thComp + rel-repr} model gives very bad fit, $\\chi^2_{\\nu} = 308/24$. Replacing the {\\sc diskbb} component with a Comptonized blackbody ({\\sc comptt}) results in a dramatic improvement, $\\chi^2_{\\nu} = 13/22$. Thus the best fit is obtained for the proper physical description of X--ray reprocessing and the {\\em Comptonized blackbody\\/} model for the soft component. The same effect is observed in GRO J1655-40, GS~2000+25 and XTE J1550-564 (\\.{Z}ycki et al., in preparation).  \\subsection{The High State}  \\begin{figure*} % \\hfil\\parbox{0.49\\textwidth}{ \\epsfysize = 4. cm \\epsfbox[50 320 550 720]{zycki_TIII-06_fig1a.ps} }\\hfil \\parbox{0.49\\textwidth}{ \\epsfysize = 4. cm \\epsfbox[50 400 550 750]{zycki_TIII-06_fig1b.ps} }\\hfil \\caption[]{Spectrum of Nova Muscae 1991 in HS, showing a soft thermal component broader than a disk blackbody.} \\label{zycki_TIII-06:fig1} \\end{figure*}  Nova Muscae 1991 was in the High State when observed on Mar 18th. The spectrum shows a strong soft component, and the hard component is negligible  \\cite{Ebisawa}. Modeling the spectrum with the {\\sc diskbb + thComp } model gives $\\chi^2_{\\nu} = 27/24$ (see Fig.~\\ref{zycki_TIII-06:fig1}). The Comptonized component is steep and makes a significant contribution to the soft component. Again, the soft component has to be described by a model {\\em broader\\/} than the multi-color disk blackbody.  \\subsection{The Very High State}  On Jan 11th, while still on the rise to the peak of the outburst, Nova Muscae 1991, showed a spectrum consisting of two Comptonized components: the softer one again {\\em cannot\\/} be described by the {\\sc diskbb} model. It can be modeled as Comptonization of $\\approx 0.3$ keV blackbody by a warm, $k \\Te \\approx 3$ keV, optically thick, $\\tau\\sim 7$, plasma.  \\subsection{Physical origin of the Comptonization}  The additional Comptonization seems to be a common feature of disk-dominated spectral states of accreting black holes. There are two possible scenarios to explain it. The first one is Comptonization in a ``hot skin'' on top of the accretion disk. Such a hot layer could not be heated by X--ray irradiation (as postulated for the low/hard state), but it would require enhanced viscous energy dissipation compared to the bulk of the disk. The second possibility is that the electron energy distribution in the Comptonizing plasma is a hybrid, thermal--non-thermal one, rather than a pure thermal (Maxwellian) distribution. In fact, the {\\sc compPS} model (J. Poutanen, private communication) using the hybrid distribution gives very good fits to  the HS and IS data sets: $\\chi^2_{\\nu} = 17/24$ for IS, $\\chi^2_{\\nu} = 16/23$ for HS, but it fails for VHS. ",
        "conclusions": "Interactions between accretion disk and corona in the soft spectral states of accreting black holes are clearly complex. One outcome of those interactions is a presence of warm, mildly optically thick plasma modifying the disk thermal emission. Proper modeling of the resulting spectra is important e.g.\\ for  correct determination of system parameters.  The 'flip-flop' transitions observed in VHS of GS~1124-68 close to the peak of its outburst suggest that the optically thick accretion flow can proceed in two modes giving the same time-integrated characteristics. One mode, usually observed in HS, shows  the $f^{-1}$-like PDS, i.e.\\ a scale-invariant behaviour, perhaps realized in magnetic turbulence \\cite{Kawaguchi}. The other mode  clearly possesses certain spatial/temporal scale, since the PDS is described by a Lorenzian, i.e.\\ it corresponds to exponential shots of a single time scale."
    },
    "0011/astro-ph0011393_arXiv.txt": {
        "abstract": "We present results of two simulations of the convection zone, obtained by solving the full hydrodynamic equations in a section of a spherical shell. The first simulation has cylindrical rotation contours (parallel to the rotation axis) and a strong meridional circulation, which traverses the entire depth. The second simulation has isorotation contours about mid-way between cylinders and  cones, and a weak meridional circulation, concentrated  in the uppermost part of the shell.    We show that the solar differential rotation is directly related to a latitudinal entropy gradient, which pervades into the deep layers of the convection zone. We also  offer an explanation of the angular velocity shear found at low latitudes near the top. A non-zero correlation between radial and zonal velocity fluctuations produces a significant Reynolds stress in that  region. This constitutes a net transport of angular  momentum inwards, which causes a slight modification of the overall structure of the differential rotation near the top. In essence, the {\\it thermodynamics controls the dynamics through the Taylor-Proudman momentum balance}. The Reynolds stresses only become significant in the surface layers, where they generate a weak meridional circulation and an angular velocity `bump'. ",
        "introduction": "\\subsection{Observations} In the outer 28 $\\%$ by radius of the Sun, known as the Solar Convection Zone (SCZ), most of the vertical energy transport is by convection. Due to the combination of convection and rotation, the gaseous body exhibits differential rotation i.e. non-uniform angular velocity. Helioseismology observations (Scherrer et al 1995, Libbrecht 1989) of the nearly 10 million acoustic p-modes that leak from the interior  into the atmosphere, have provided through frequency splittings, information on the angular velocity distribution of the solar interior. The latest results suggest that  the isorotation surfaces in the solar convection zone are cone-like (aligned radially), in disagreement with most numerical simulations, which tend to produce cylindrical contours parallel to the rotation axis. Specifically, such observations have revealed that: the angular velocity near the equator first increases then gently decreases with depth, at mid-latitudes it is almost constant with depth, while at high latitudes it increases with depth; on the surface of the Sun, the angular velocity increases from the poles (35 day period) to the equator (25 day period) and  meridional circulation $v_\\theta$,  is much weaker than differential rotation (zonal velocity  $\\approx  2{\\rm km/s}, v_\\theta \\approx 25 {\\rm m/s}$).   \\subsection{Simulations} \\subsubsection{Computational hurdles} Modeling the SCZ is a  formidable task. Observations suggest: that the SCZ is highly stratified, spanning about 19 pressure scale heights in depth, with the Mach number (square of the ratio of flow velocity to sound speed)  approaching unity at the top; the Prandtl number (ratio of the timescales of thermal to viscous diffusion) is extremely small ($10^{-6}$); the equation of state is complex; and finally, the motion is highly turbulent, revealed by the large number of scales ranging  from 0.1 km to $10^5$ km (depth of SCZ). In the photosphere (near the top of the SCZ) the kinematic viscosity of hydrogen is about $0.1{\\rm cm}^2{\\rm s}^{-1}$, while the root mean square (r.m.s.) velocity in the SCZ is the order of $100 {\\rm m}{\\rm s}^{-1}$. These values suggest a  Reynolds number, ${\\rm Re}$=velocity $\\times$ length / kinematic viscosity, of at least $10^{12}$. As the number of degrees of freedom needed to represent a flow is proportional to ${\\rm Re}^{9/4}$, to resolve numerically the scales in all three directions  would require $10^{27}$ grid points. Even with present technology the maximum Reynolds number modeled in direct numerical simulations (DNS) is a few thousand (Brummel, Hurlbert \\& Toomre 1996). Alternatively, if one assumes most of the energy is contained in the resolved scales, which are much larger than the viscous and thermal dissipation scales (unresolved scales), and that the average energy transferred by these smaller scales can be modeled by entropy diffusion, then a much larger viscosity (smaller ${\\rm Re}$) can be used to model the turbulent flow. This is the principle of large eddy simulations (LES).   \\subsubsection{Numerical simulations  of the SCZ} Explicit integration of the hydrodynamic equations requires the timestep to be less than the time for a signal (sound wave) to travel between two grid points. This is known the Courant-Friedrichs-Levy (CFL) stability criterion. Previous investigations have solved approximate forms of the fully compressible equations, in which  sound waves do not  exist. This circumvents  the use of restrictively small timesteps. Incompressible convection in a shell was studied by Gilman (1978)  using   the Boussinesq approximation.  This  allows  density variation only when coupled with gravity. A better approximation is to exclude only temporal variations in density, known as the anelastic approximation. This still suppresses sound waves but allows density to vary vertically. The most recent anelastic simulation in a full shell (global model) on massively parallel architectures are by Miesch et al (2000). Though there were significant improvements on earlier work (Glatzmaier 1987), they were still unable to adequately resolve the observed cone-like structure of the angular velocity contours. Overall the isorotation contours were cylindrical. Furthermore, the angular velocity shear layer near the top of the convection zone could not be resolved. Near the top, vigorous compressible motions have Mach numbers close to one and the anelastic approximation breaks down.  The fully compressible equations have however been solved in local models (boxes). By considering only a tiny section of a spherical shell, Brummel et al. (1996)  did a DNS of turbulent compressible convection in f-planes.  In essentially the same geometry, Chan (1999)  did a set of LES computations, in f-planes.  The LES model emphasized efficient convection (convective flux $\\gg$ diffusive flux) as opposed to the inefficient convection in the DNS models. In f-plane simulations the angle between the rotation vector and gravity is considered constant throughout the box. The  problem with  these local models is that they have periodic horizontal boundaries. In such domains the only possible source  of differential rotation is the Reynolds stress. Horizontal averaging removes latitudinal gradients of thermal quantities  which are a possible source of the differential rotation. For example, in the model by Durney (1999), the latitudinal entropy gradient (or baroclinicity) is essential in shaping the isorotation contours. Additionally, the f-plane models do not allow a realistic  meridional circulation to develop; nor can Rossby waves exist as the rotation vector is constant throughout the domain. Rossby waves have been  suggested as a possible source of correlation between longitudinal and meridional motions (Brummel et al. 1996).  Our model represents a compromise between the anelastic global and compressible local simulations. We solve the full Navier Stokes  equations in  spherical coordinates in a domain with a significant latitudinal coverage. In this way we are able to resolve the upper shear layers and  simultaneously include the effects of meridional circulation and latitudinal gradients of thermal quantities. To circumvent the CFL restriction on the time step, semi-implicit time integration is employed. We use the Alternating Direction Implicit Method on a Staggered Mesh (ADISM) (Chan \\& Wolff 1982)   in conjunction with an explicit method. Due to computational restraints of three-dimensional calculations, we are limited to studying a small section of the entire shell. A full simulation on a  $70 \\times 70 \\times 39$ three-dimensional mesh, spans just $60^\\circ$ in longitude and latitude. This calculation requires about a month to run when parallelized on 4 processors of the ORIGIN 2000.  This paper consists of 4 more sections. Section 2 describes the overall physical setup, formulates the mathematical model and describes the numerical approach. This is followed by descriptions  of the statistically averaged zonal and meridional flows, and  turbulent nature of the compressible convection. The third section attempts to pin down the source of the differential rotation and meridional circulation. The final section is a conclusion. ",
        "conclusions": "The aim of this work is to reproduce solar differential rotation, by solving the equations of hydrodynamics, in a section of a spherical shell. To achieve this, we have done two numerical simulations. The first simulation initially has  an angular velocity that  decreases inwards, in agreement with the solar case. However, as the computation progresses, the radial angular velocity gradient changes sign, and the rotation rate increases inwards. The mean rotational structure consists of cylindrical isorotation contours and a sub-rotating equator. The switch is due to an artificially large meridional circulation, which itself is a result of strong downflows that occur at the impenetrable side boundaries.  In the second computation the longitudinal span is doubled, so that it now equals the latitudinal span. The wider shell has much weaker downflows and a more turbulent flow. Under these  conditions the rotation profile remains in the `sun-like' state.  The angular velocity now bears a closer resemblance to solar case, with a radial and latitudinal variation, both qualitatively and quantitatively similar to the SCZ.  By using an implicit timesteping scheme,  we are able to run the simulations for longer than a `Kelvin Helmholtz' time scale. The emphasis on complete thermal relaxation, could be the vital ingredient, required to correctly model rotating convection. Incomplete thermal relaxation, may be one reason why earlier simulations, failed to get the proper rotation pattern. In the `Sun-like' simulation,  it is the {\\it large}  scale thermal structure (specifically the entropy variation with latitude),  rather than {\\it small} scale motions (as in mean  field models or some numerical models), that directly produces the differential rotation. The non-zero latitudinal entropy gradient (baroclinicity) shapes the differential rotation. We notice that in the `sun-like' model, the kinetic energy distributes roughly equally between the mean and turbulent scales, in contrast to the non `sun-like' model, in which nearly all of the energy is in the mean flow. Furthermore, the turbulent kinetic energy is greatest in the zonal direction. This suggests the turbulent nature of the SCZ, may have some indirect role in the maintenance of solar differential rotation, but how operates is still unclear.  The  Reynolds stresses are important in the upper 10 \\% of the computational domain. This region contains the top of the convection (unstable) layer and all of the radiative (stable) layer. Here the velocity correlation $\\overline{v_r' v_\\phi'}$, has a dual effect. Firstly, it provides an inward transport of angular momentum, which causes the slight increase in angular velocity, just below the surface and  near the equator. Secondly, $-\\pa / \\pa r (\\overline{v_r' v_\\phi'})$ is principally responsible for the poleward meridional circulation at the top of the convection layer, and the accompanying return flow in the stable layer. Closer to the equator (within about $5^\\circ$), $-\\pa / \\pa \\theta (\\overline{v_\\theta' v_\\phi'})$ appears to drive the meridional flow. In the meridional momentum equation, the equator excluded,  the only significant terms are  the pressure gradient and the Coriolois force. The Reynolds stresses are only important in  the  zonal momentum equation, here they generate the meridional circulation found at the top of the shell.  The key result is that the dynamics near the top (surface layers) of the convection zone,  are controlled by the Reynolds stresses, while elsewhere, the differential rotation is determined the Taylor-Proudman momentum balance. A similar conclusion was reached in the semi-analytical model of the SCZ by Durney (2000) (and references therein)."
    },
    "0011/astro-ph0011446_arXiv.txt": {
        "abstract": "Thermal and chemical evolution of primordial gas clouds irradiated with far-ultraviolet (FUV; $h \\nu < 13.6$ eV) radiation is investigated. In clouds irradiated by intense FUV radiation, sufficient hydrogen molecules to be important for cooling are never formed. However, even without molecular hydrogen, if the clouds are massive enough, they start collapsing via atomic hydrogen line cooling. Such clouds continue to collapse almost isothermally owing to successive cooling by H$^{-}$ free-bound emission up to the number density of $10^{16}{\\rm cm^{-3}}$. Inside the clouds, the Jeans mass eventually falls well below a solar mass. This indicates that hydrogen molecules are dispensable for low-mass primordial star formation, provided fragmentation of the clouds occurs at sufficiently high density. ",
        "introduction": "Since the dawn of evolutionary cosmology, great efforts toward understanding the formation of the first structures have been made by many authors (e.g., Saslaw \\& Zipoy 1967; Peebles \\& Dicke 1968; Matsuda, Sato, \\& Takeda 1969; Doroshkevich \\& Kolesnik 1976; Carlberg 1981; Palla, Salpeter, \\& Stahler 1983; Kashlinsky \\& Rees 1983). In particular, partly stimulated by observational advances and by the consequent detectability of the first stars by next-generation facilities, this area is being vigorously investigated these days (e.g., Haiman, Thoul, \\& Loeb 1996; Ostriker \\& Gnedin 1996; Tegmark et al. 1997; Yamada \\& Nishi 1998; Bromm, Coppi, \\& Larson 1999; Abel, Bryan, \\& Norman 2000; Fuller \\& Couchman 2000). Those studies have repeatedly emphasized the key role played by molecular hydrogen.  The role of H$_2$ in the cosmological context is twofold. First, small primordial clouds whose virial temperature is $1000 \\la T_{\\rm vir} \\la 8000{\\rm K}$ can cool and start collapsing via H$_2$ line cooling after virialization (e.g., Haiman et al. 1996). Second, the Jeans mass inside a collapsing primordial cloud is eventually reduced to much less than a solar mass as a result of H$_2$ cooling (e.g., Palla et al. 1983).  Regarding the former role, there is some dispute over the role of H$_2$ in the formation of second-generation stars, namely, what happens after the first stars have formed by H$_2$ cooling. Formed first stars or first quasars can be sources of UV radiation. Whereas ionizing photons are trapped in an HII region around the source, far-ultraviolet (FUV; $h\\nu <13.6$ eV) photons \\footnote{Some authors use the term ``soft-UV'', instead of ``far-UV'', for radiation whose energy is below the Lyman limit.} travel farther and form an outer H$_2$ photodissociation region (PDR). Therefore, before the overlap of the HII regions, or in other words the reionization of the universe, the H$_2$ PDRs must cover the whole intergalactic space. At that time, the universe is filled with FUV background radiation (e.g., Haiman, Rees, \\& Loeb 1997; Ciardi, Ferrara, \\& Abel 2000).  Hydrogen molecules in small objects can easily be photodissociated by either internal or external UV sources. Haiman et al. (1997) found that in the presence of UV background radiation at the level needed to reionize the universe, molecular hydrogen is photodissociated in small cosmological objects. Furthermore, Omukai \\& Nishi (1999) pointed out that even a single O star formed in a small cosmological object is able to photodissociate the whole original cloud. In addition, those clouds are easily blown out by a few supernovae because of their shallow potential wells (e.g., Mac Low \\& Ferrara 1999). These effects strongly regulate the subsequent star formation within small pregalactic clouds.  In contrast to this, the temperature of sufficiently massive clouds reaches above $\\sim$8000K, where atomic Ly$\\alpha$ emission begins to be an effective cooling agent. Regardless of the H$_2$ photodissociation, such clouds can start collapsing dynamically. It is frequently assumed that, even in primordial clouds that cool by the atomic cooling instead of the H$_2$ one, star formation occurs in the same way as in those with the H$_2$ cooling. However, it is far from evident.  Recall another role of hydrogen molecules in primordial star formation. The formation of protostars in primordial clouds with H$_2$ cooling has been investigated by many authors (e.g., Matsuda et al. 1969; Yoneyama 1972; Carlberg 1981; Palla et al. 1983; Omukai \\& Nishi 1998). It is known that those clouds cool by H$_2$ rovibrational line emission (for number density $n \\la 10^{14} {\\rm cm^{-3}}$), H$_2$ collision-induced emission (for $10^{14} \\la n \\la 10^{16} {\\rm cm^{-3}}$), and the H$_2$ dissociation (for $n \\la 10^{20}{\\rm cm^{-3}}$). Eventually, a very small protostar ($\\sim 10^{-3} M_{\\sun}$) forms at the center when $n \\simeq 10^{22}{\\rm cm^{-3}}$.  On the other hand, we know little about the later evolution of clouds that start collapsing via atomic cooling. It is still unclear under what condition H$_2$ eventually forms, or whether it plays some role as a coolant in such clouds. It is also not known how small a Jeans mass is reached through atomic cooling alone. If this value falls well beyond the mass of normal stars (for example if it corresponds to the size of super massive stars or very massive objects), we must conclude that normal stars does not form only by atomic cooling. The aim of this paper is to answer these questions.  In this paper, we study the thermal and chemical evolution of primordial clouds irradiated with FUV radiation. Thereby, we demonstrate that the Jeans mass inside the clouds is reduced to much less than a solar mass even in the absence of molecular cooling. This implies that low-mass primordial star formation is possible even without H$_2$ if fragmentation of the clouds occurs at sufficiently high density.  The outline of this paper is as follows. In \\S 2, the method of our calculations is described. In \\S 3, results of our calculations are presented. Finally, our work is summarized and its implications sketched briefly in \\S 4. ",
        "conclusions": "We have investigated the thermal and chemical evolution of primordial clouds irradiated with FUV radiation. When intensity of the irradiating FUV radiation exceeds some critical value [$J_{21} \\simeq 10^{5}$ for $J_{\\rm UV}(\\nu) \\propto \\nu^{-1}~~( h \\nu < 13.6 {\\rm eV})$], sufficient molecular hydrogen to be important for cooling is never formed in those clouds because of the photodissociation and blocking of the formation channel. Nonetheless, sufficiently massive clouds can start dynamical collapse by atomic line cooling even without H$_2$. Those clouds continue to collapse almost isothermally at several thousand K along the ``atomic cooling track'' in the density-temperature plane. That dynamical collapse is induced by successive cooling by Ly$\\alpha$ emission, two-photon emission, and H$^{-}$ free-bound emission. The minimum Jeans mass eventually reduces to about $0.03 M_{\\sun}$ for those clouds. On the other hand, clouds irradiated with less intense FUV radiation collapse dynamically by the H$_2$ cooling. The minimum Jeans mass of such clouds is about $1.5 \\times 10^{-3} M_{\\sun}$.  Although in this paper we limited our application to the protostellar collapse under an FUV background, we expect that the primordial protostellar collapse without H$_2$ also occurs in other situations. Recently, Susa \\& Kitayama (2000) investigated the collapse of primordial clouds irradiated with ionizing UV background $J_{\\rm UV}(\\nu)=J_{21} \\times 10^{-21}(\\nu/\\nu_{\\rm th})^{-1}~{\\rm (erg~s^{-1} cm^{-2} str^{-1} Hz^{-1})}$. They pointed out that H$_2$ formation is suppressed for $J_{21}>3.6 \\times 10^{4}$ owing to blocking of the H$_2$ formation channel by H$^{-}$ photodissociation. Although they did not study the later evolution, we guess that the clouds continue collapsing along our atomic cooling track. Carlberg (1981) also considered the thermal evolution of collapsing primordial clouds without external radiation, but he did not include the three-body process of ${\\rm H_2}$ formation. For this reason, in his calculation, molecular hydrogen dissociates at $n \\simeq 10^{11} {\\rm cm}^{-3}$. After that, the evolutionary path of his cloud follows approximately our atomic cooling track (see Fig. 2 of Carlberg 1981). This fact support our speculation that the clouds that do not cool by H$_2$ collapse along the atomic cooling track.  Next, we discuss the feasibility of the collapse along the atomic cooling track in the cosmological reionization. The UV intensity needed to reionize the universe at $z_{\\rm reion}$, i.e., 1 ionizing photon per a hydrogen nucleus, is \\begin{equation} J_{\\rm reion}=\\frac{hc \\Omega_{\\rm b} \\rho_{\\rm cr} (1+z_{\\rm reion})^{3}}{4 \\pi (1+4 y_{\\rm He}) m_{\\rm H}}=3 \\times 10^{-21} (\\frac{1+z_{\\rm reion}}{10})^{3} ~{\\rm (erg~s^{-1}~cm^{-2}~str^{-1}~Hz^{-1})} \\end{equation} at the Lyman limit for $J(\\nu) \\propto \\nu^{-1}$ and the assumed cosmological parameters. Thus, the condition that the clouds collapse along the atomic cooling track, $J_{21} \\ga 10^{5}$ for $J_{\\rm UV}(\\nu) \\propto \\nu^{-1}$, seems to be rarely met in practice, except in the immediate vicinity of UV sources.  Finally, we discuss the mass of formed stars. Although the minimum Jeans mass is less than 0.1 $M_{\\sun}$, this corresponds not to the final mass of formed stars, but rather to the initial mass of the protostar, which has no direct relation with the former. When a small hydrostatic core, namely protostar, forms at the center of a collapsing protostellar cloud, an enormous amount of gravitationally unstable gas still surrounds the protostar. Thereafter, the protostar grows in mass by accreting the ambient matter. Consequently, the final mass of stars is determined by the subsequent accretion onto the protostar. Stahler, Shu, \\& Taam (1980) argued that a rough estimate of the protostellar mass accretion rate ${\\dot M}$ can be obtained from the relation ${\\dot M} \\sim c_{\\rm s}^{3}/G$, where $c_{\\rm s}$ is the isothermal sound speed in the protostellar cloud. From this relation, the mass accretion rate is higher for protostars formed by atomic cooling than for those formed by molecular cooling because of the higher temperature of the protostellar cloud. Suppose that the accretion is halted by some stellar activity, e.g., radiation force, bipolar flow, etc. In this case, the higher accretion rate probably results in the later halting and higher mass of formed stars (e.g., Wolfire \\& Cassinelli 1987).  In contrast, from the viewpoint of fragmentation, the masses of primordial stars formed by atomic cooling alone could be smaller than those formed by molecular cooling. The fragmentation is possible in principle as long as the temperature decreases toward higher density, while it does not occur after the isothermality of the cloud breaks down and the temperature begins to rise (Inutsuka \\& Miyama 1997; Masunaga \\& Inutsuka 1999). From Figures 2 and 3, the final fragmentation occurs at about $n \\simeq 10^{16} {\\rm cm^{-3}}$ for clouds collapsing along the atomic cooling track. On the other hand, it seems difficult for molecular cooling clouds to fragment at $n \\ga 10^{3} {\\rm cm^{-3}}$. The masses of fragments, which is about the Jeans mass at that time, are about $0.1 M_{\\sun}$ for the atomic-cooling clouds, and $10^{4} M_{\\sun}$ for the molecular-cooling clouds. If this is the case and the whole material contained in a fragment is converted into a single star, very massive stars form from a molecular-cooling cloud, and red dwarfs form as a result of atomic cooling. However, it should be noted that the fragmentation and thermal evolution depend strongly on each other and also on the geometry of the clouds. In fact, Uehara et al. (1996) concluded that the minimum mass of fragments is on the order of a solar mass, which is essentially the Chandrasekhar mass, by studying thermal evolution of filamentary primordial clouds that collapse by molecular cooling. Later, improved simulations by Nakamura \\& Umemura (1998) confirmed Uehara et al.(1996)'s result. The discrepancy between our mass of fragments ($\\sim 10^{4} M_{\\sun}$) and theirs ($\\sim 1 M_{\\sun}$) results from the difference of the assumed geometry of the clouds; we have assumed spherical clouds in studying the thermal evolution, while their clouds are filamentary. For filamentary clouds, the gravitational contraction is slower and compressional heating rate is lower than for spherical ones. Hence, for the filamentary cloud the isothermality breaks down at higher density, where the Jeans mass is smaller. To fully address the complexity arising from the cloud geometry, two-dimensional or three-dimensional calculations are needed. Furthermore, there are other uncertainties relating to the evolution of the fragments, for example, mass accretion onto the fragments, merger between them, feedback from neighboring formed stars, etc. (e.g., Bromm et al. 1999). It seems premature to draw conclusions here about these issues."
    },
    "0011/astro-ph0011500_arXiv.txt": {
        "abstract": "We present a simple method for adaptively binning the pixels in an image. The algorithm groups pixels into bins of size such that the fractional error on the photon count in a bin is less than or equal to a threshold value, and the size of the bin is as small as possible. The process is particularly useful for generating surface brightness and colour maps, with clearly defined error maps, from images with a large dynamic range of counts, for example X-ray images of galaxy clusters.  We demonstrate the method in application to data from \\emph{Chandra} ACIS-S and ACIS-I observations of the Perseus cluster of galaxies. We use the algorithm to create intensity maps, and colour images which show the relative X-ray intensities in different bands. The colour maps can later be converted, through spectral models, into maps of physical parameters, such as temperature, column density, etc. The adaptive binning algorithm is applicable to a wide range of data, from observations or numerical simulations, and is not limited to two-dimensional data. ",
        "introduction": "A problem often encountered with the processing of images, especially those from X-ray telescopes, is that there is a large dynamic range present in intensity. In order to examine the structure of emission in regions of an image where there are few photons, it is necessary to bin the data using bins of large angular size. An alternative is to smooth the data with a Gaussian kernel of large angular size. Conversely, to examine structure in areas of high emission, a small bin-size or kernel is sufficient. If one wishes to examine regions with a wide variation in count-rate, then some procedure which has a variable bin or kernel is required to create statistically optimal maps.  Another related problem, which motivated the present work, is the creation of optimal colour maps.  The ratio of photon counts in a low intensity region will have a large associated statistical error, so the counts need to be averaged over a large region to form the ratio. The algorithm needs to account for the statistical accuracy in the separate bands, not just the total image.  One technique which allows structure on a wide range of scales to be revealed is adaptive kernel smoothing (Ebeling, White \\& Rangarajan 2000; Huang \\& Sarazin 1996). The \\textsc{asmooth} algorithm of Ebeling et al. (2000) convolves an image with a Gaussian kernel. If the signal within the kernel, applied to part of an image, is of a chosen minimum significance above the local background, then the convolved signal is added to the output image. The kernel is increased in size until all the counts in the input image are added to the output image.  \\textsc{asmooth} is very good at finding low surface brightness features, such as filaments. It also especially useful for producing images for display, and for identifying bright features against a background.  It is not a universal smoothing tool, however, and care needs to be taken before using it in a particular situation. For instance, it is designed to identify positive features against a background, but not to identify `holes' in emission. Due to the way it identifies significance against the local background, it may produce spurious features in areas where there is bright flat emission, for instance the cores of galaxy clusters, where Poisson noise may become significant.  Used with care, \\textsc{asmooth} is a useful routine in many situations.  The analysis method outlined in this paper, \\emph{adaptive binning}, was originally developed to produce temperature, column density and colour maps of the Perseus cluster in Fabian et al. (2000). We required a simple algorithm to compare the relative fluxes in different X-ray bands to a theoretical plasma emission model. In order to be able to spatially resolve the ratios, it is necessary to bin or smooth the data.  Binning, rather than smoothing, is a better method for producing spatial estimates of counts for analysis, because it does not spread counts around an image beyond the boundaries of the bin.  In this paper we use data analysed from the 24\\,ks ACIS-S \\emph{Chandra} X-ray observation of the Perseus Cluster, Abell\\,426, from Fabian et al. (2000), to demonstrate the adaptive binning method. The Perseus Cluster is the brightest X-ray cluster in the sky, and contains a large cooling flow (Fabian et al. 1994) with a mass deposition rate of approximately $300\\Msunpyr$. Interesting substructure in the cluster due to its central radio source makes it an ideal object for our purpose.  We will not attempt here to analyse the physics behind the observations, which is left to Fabian et al. (2000) and further work. We also study the method by applying it to simulated data.  The algorithm is presented in a logically consistent way, from intensity adaptive binning to colour adaptive binning. Colour binning, as mentioned above, was our original motivation, developed to analyse X-ray colour cluster images. ",
        "conclusions": "We have presented an algorithm to adaptively bin data. We have demonstrated its use on X-ray intensity and colour images of the Perseus cluster. We also applied the algorithm to simulated data. The method has also been used to create multi-layer colour images, where the colours represent the intensity in different bands. The fractional error threshold can be varied depending on how much statistical noise is acceptable in the analysed image.  The method is not limited to binning on intensity or colour. Any sort of data may be binned, providing there is a method to compute the statistical fractional error, or weighting, of a group of pixels. The algorithm may be useful in presenting the output from numerical simulations.  The method is also not limited to two-dimensional data. We have used it already on one-dimensional cuts through images. It may also be simply extended to work on $N$-dimensional data sets."
    },
    "0011/astro-ph0011514_arXiv.txt": {
        "abstract": "Multiple systems play an important role in the evolution of star clusters. First we discuss several formation mechanisms which depend on the presence of binaries, either primordial or of dynamical origin. Hierarchical configurations are often stable over long times and yet may experience evolution of the internal orbital parameters. We describe an attempt to model the eccentricity change induced by the outer component using an averaging method, together with the effects due to tidal dissipation and apsidal motion acting on the inner binary. This treatment is adopted for systems with high induced eccentricity which gives rise to some interesting outcomes of significant period shrinkage. ",
        "introduction": "$N$-body simulations of star clusters need to include a significant fraction of primordial binaries in order to be realistic (Raboud \\& Mermilliod 1999). This introduces a new complexity into the problem and also poses many technical challenges. Among the main computational aspects we mention strong interactions and short time-scales, as well as astrophysical phenomena involving tidal interactions or mass transfer in short-period binaries. In the present paper, we are concerned with the formation and evolution of multiple systems containing different hierarchical configurations. Let us introduce the notation B and S for binaries and single stars. In addition to standard triples, denoted by [B,S], we distinguish quadruples in the form [B,B] as well as higher-order systems, such as [[B,S],S] up to sextuplets [[B,B],B] which are all recognized by the algorithms.  It is known from scattering experiments that hierarchical triples may form as the result of binary-binary collisions, where the fourth component is ejected (Mikkola 1983). Provided the pericentre of the outer orbit is sufficiently wide, such configurations may be stable over long times (Mardling \\& Aarseth 1999). Hence it only requires a small formation rate in order to build up a substantial population of such systems. Given long life-times, more complex structures may be formed as the result of favourable interactions. In the present contribution, we consider the formation process in more detail by examining some numerical results and also model the inner binary evolution which is subject to eccentricity modulations. ",
        "conclusions": "The first part of this paper deals with the formation of hierarchical systems in star cluster models containing a significant proportion of primordial binaries. Provided the orbital characteristics satisfy the stability criterion given by Eq.~(1) and the binding energy of the outer component is fairly hard, a triple is represented by the corresponding two-body motion. We find that typically at least ten such systems are present after the first $Gyr$ in simulations with $N_s = 8000$ single and $N_b = 2000$ binary stars. The stable triples form as the result of binary-binary collisions, often resulting in hard outer energies and are therefore quite robust. Moreover, the relatively small formation rate is compensated by long life-times. On the other hand, most quadruples originate via classical three- and multi-body processes and are therefore less well bound initially. Given the presence of such systems, further interactions may produce more complex structures like quintuplets and sextuplets, especially during late stages of the cluster life-time. Thus we may invoke the concept of dynamical molecules when describing the different outcomes of interactions (cf.\\ Aarseth 1988). Consequently, the internal degrees of freedom of high-order systems represent latent energy which becomes available for heating the cluster on disruption. We also emphasize that hierarchical systems have large cross sections and tend to be centrally concentrated. Since such configurations originate by a combination of processes, their formation is best studied {\\it in situ}, rather than by scattering experiments which have so far only reached the case $B + B$ in complexity.  In the second part, we outline a scheme for modelling the eccentricity change of the inner binary due to the outer component moving in an inclined orbit. Thus if the inclination is favourable, the induced eccentricity may attain a large value, with a corresponding small pericentre separation, and thereby initiate tidal dissipation. In some cases this leads to significant orbital shrinkage at the end of the circularization and is therefore a mechanism for producing close binaries. Our development is based on calculating the effect of the outer body on the inner eccentricity and angular momentum vectors, using a quadrupole expansion. In addition, we include contributions from the internal tidal dissipation and apsidal motion due to oblateness, as well as relativistic precession. So far the averaging treatment appears to be promising but should still be considered experimental."
    },
    "0011/astro-ph0011088.txt": {
        "abstract": "The kinematic and Period--Luminosity--Colour distribution of O-rich Long-Period Variable stars of the solar neighbourhood is interpreted in terms of pulsation modes, masses and metallicities. It is first shown that, because of input physics imperfections, the periods and mean colours derived from  the existing linear and nonlinear nonadiabatic models must significantly  depart from the actual behaviour of the stars. As a consequence systematic corrections have to be applied, as a first approximation, to our linear  model grid. These free parameters, as well as the mixing length, are  calibrated on the LPVs of the LMC and of some globular clusters, assuming  a mean mass of $1\\,M_\\odot$ for the LMC Mira-like stars. Then, the  masses and metallicities corresponding to the four kinematic/photometric  populations of local LPVs are evaluated. The possibility of a varying   mixing-length parameter is discussed and taken into account. Stars of the old  disk appear pulsating in the fundamental mode\\,: one group, mainly composed  of Miras, has mean mass $<M>\\;\\simeq 0.9\\,M_\\odot$ and mean metallicity  $<Z>\\;\\simeq 0.02$, both strongly increasing with the period\\,; a second  group, slightly older and mainly composed of SRb's, has  $<M>\\;\\simeq 0.9\\,M_\\odot$ and $<Z>\\;\\ga 0.03$\\,. Stars of the  thin disk appear pulsating in the first and second overtones, with  $<M_{1ov}>\\;\\approx 1.05\\,M_\\odot$, $<M_{2ov}>\\;>0.75\\,M_\\odot$ and  $<Z>\\;\\ge 0.04$\\,. Stars of the extended disk/halo appear pulsating in  the fundamental mode, with $<M>\\;\\simeq 1.1\\,M_\\odot$ and $<Z>\\;\\simeq 0.01$. The mixing-length parameter probably decreases along the AGB by no more than 15\\,\\% per magnitude. The large, positive period corrections (more than 30\\,\\% for the fundamental and 8\\,\\% for the  first overtone) that have to be applied to the LNA models used in this  study do not seem to be explained by imperfect sub-photospheric physics alone,  especially when nonlinear effects are taken into account. The origin of the  extra period increase (at least 15\\,\\% for the fundamental mode) may be the  stellar wind, which was neglected by all pulsation codes up to now. ",
        "introduction": "Among the most important and still unsolved issues concerning Long-Period Variable stars (LPV) are the modelling of their pulsation and even the mere identity of the predominant pulsation mode. Linear nonadiabatic (LNA) models are currently used on a wide scale for reasons of simplicity. Relying on LNA relations between the period and the fundamental parameters, and on dynamical models of the photospheric region, the pulsation mode of Mira stars is the fundamental according to the observed velocity amplitude (Hill \\& Willson 1979, Willson 1982, Bowen 1988, Wood 1990), and the first overtone according to the angular diameter estimates (Haniff et al. 1995, Van Leeuwen et al. 1997), pulsation accelerations (Tuchman 1991) and the PL distribution of globular cluster stars (Feast 1996). One study based on the periods and luminosities of LPVs of the Large Magellanic Cloud (LMC) concluded that Miras were pulsating in the fundamental mode and low-amplitude semi-regulars on the first or second overtone (Wood \\& Sebo 1996). However, another study, also taking into account a few Miras of the solar neighbourhood and various temperature estimates, supported the first overtone in a majority of Miras (Barth\\`es 1998). Last, still relying on LNA models, Barth\\`es \\& Tuchman (1994) and Barth\\`es \\& Mattei (1997) found that the Fourier spectra of a few Miras were more easily explained than were the stars pulsating on the first overtone with the fundamental mode and other overtones also acting.  In fact, the large pulsation amplitude and the complexity of the chemistry and radiative transfer in the photospheric region make the mean effective temperature and radius very uncertain, even for nearby stars (Bessell et al. 1989a, 1996\\,; Hofmann et al. 1998\\,; Ya'ari \\& Tuchman 1998). The {\\sc Hipparcos} parallaxes, too, usually have large error bars, and they concern only about two dozen Miras.  Last, most of these studies suffered from a strong sampling bias, favouring periods longer than 250 days, especially those ranging from 300 to 400 days.\\\\  On the other hand, hydrodynamical models predict that, after thermal relaxation, Miras pulsate in the fundamental mode with a period either much shorter than (Ya'ari \\& Tuchman 1996, 1999) or very close to (Wood 1995, Bessell et al. 1996, Hofmann et al. 1998) the LNA period. These models mainly differ by their handling of the convective energy transport and of the equation of state, but they share the same unrealistic assumption\\,: no wind and  no extended circumstellar envelope at the outer boundary. As a consequence,  the reliability of both the nonlinear and linear nonadiabatic pulsation  models is uncertain, and the pulsation mode is still ambiguous.\\\\  An unfortunate result of these theoretical difficulties is that the present masses and metallicities of the LPVs have always been very uncertain. Indeed, no direct estimate of mass in binary systems has been possible up to now, and the chemistry and radiative transfer in the photospheric region are so complex that no reliable metallicity has been derived.\\\\  This paper mainly consists in exploiting the classification (based on kinematic and photometric criteria) and the period--luminosity--colour (PLC) distributions of the Oxygen-rich LPVs of the solar neighbourhood, which were determined in Paper~I of this series (Barth\\`es et al. 1999). Its aims are to assess the pulsation models, and to identify the predominant mode and estimate the average mass and metallicity for each kinematic/photometric group.  The data sets are presented in Sect.~2. They include the results of Paper~I, i.e. the periods, absorption-corrected absolute magnitudes and de-reddened colours of a sample of stars observed by {\\sc Hipparcos}, but also the de-biased distributions of the four kinematic/photometric groups to which they belong. LPVs found in the LMC and in various globular clusters are also included, with a view to calibrating the free parameters of the models.  Sect. 3 describes the pulsation models to which these data will be confronted\\,: the linear nonadiabatic modelling code and the adopted temperature scale, i.e. colour--temperature--metallicity (CTZ) relations. Sect. 4 then explains the difficulties to be expected because of the physical approximations of the models, and also because of the nonlinear shape of the CTZ relations. We will be led to introduce some systematic correction parameters for the periods and colours.  These free parameters, together with the mixing length, are calibrated in Sect. 5 by confronting the models with the PLC distributions of the LPVs in the LMC and globular clusters. The so-calibrated models are then confronted with the PLC distribution of local stars in Sect. 6, and the pulsation mode, mass  and metallicity are derived for each kinematic/photometric group. Finally,  the stability and reliability of the results, as well as their consequences  concerning the existing pulsation codes, are assessed in Sect.~7. ",
        "conclusions": "The aim of this paper was to interpret the results of Barth\\`es et al. (1999) in terms of pulsation modes and  fundamental parameters,  i.e. the  fact that the Long-Period Variable stars of the solar neighbourhood are  distributed among four groups (according to kinematic and photometric criteria),  and to study the period--luminosity--colour distributions of these groups.  This was done by confronting them with a grid of linear nonadiabatic pulsation  models.  Preliminary discussion of the colour--temperature relations and of the existing linear and nonlinear modelling codes showed that the periods and colours predicted by LNA models may significantly differ from the observed values, even if the adopted fundamental parameters are the true ones.  In order to mimic this behaviour, we added a few free  parameters to the LNA models: a systematic period correction $\\Delta\\log P$ for each mode, and  systematic colour corrections $\\Delta(V-K)$ and $\\Delta(J-K)$. These  parameters, as well as the mixing length, were then calibrated by demanding  that consistent masses be derived from the period and from the colour, for  the fundamental and first-overtone pulsators of the LMC and of globular  clusters with LMC or SMC metallicity. It was assumed that the mean mass of  the Mira-like stars (fundamental pulsators) of the LMC is 1 $M_\\odot$. Then,  the mean mass of the LMC first-overtone pulsators was found to be about 0.95  $M_\\odot$, while that of the fundamental and first-overtone pulsators with  SMC metallicity was 0.8 and 0.6 $M_\\odot$ respectively (or down to 0.75 and  0.5 $M_\\odot$ if the mixing-length parameter is allowed to strongly vary with  the luminosity).  We were thus able to determine the pulsating mode and the mean masses and metallicities of the neighbouring LPV populations\\,:  \\begin{itemize} \\item Group~1 stars (old disk LPVs, mostly Miras) are pulsating on the fundamental mode. The mean mass is $<M>\\,\\simeq\\,0.9\\,M_\\odot$ and the mean metallicity $<Z>\\,\\simeq\\,0.02$, both strongly increasing with the period. \\item Group~2 stars (old disk LPVs, mostly SRb's) are pulsating on the fundamental mode also, with $<M>\\,\\simeq\\,0.9\\,M_\\odot$ and $<Z>\\,\\ga\\,0.03$. Metallicity variations of about 50\\,\\% are likely. \\item Group~3 stars (thin-disk LPVs, mostly SRb's) are pulsating on the first and second overtones, with \\linebreak[4] \\mbox{$<M_{1ov}>$}$\\;\\approx\\,1.05\\,M_\\odot$, \\mbox{$<M_{2ov}>$}$\\;>\\,0.75\\,M_\\odot$ and \\linebreak[4] $<Z>\\;\\geq\\,0.04$. \\item Group~4 stars (extended disk and halo) appear to be pulsating on the fundamental mode with $<M>\\,\\simeq\\,1.1\\,M_\\odot$ and $Z\\,\\simeq\\,0.01$. This result must however be taken with caution. \\end{itemize}  These findings, based on the $V-K$ index as a temperature estimator, have been checked and confirmed by using $J-K$. They were proven to be stable and consistent with the kinematics. The error bars of the masses are about 5\\,\\%. If one shifts the reference mass of the LMC Miras by any reasonable amount, all other masses are simply similarly shifted and the metallicities remain nearly unchanged. Metallicity differences with respect to the LMC may be a little overestimated if molecular opacities have been significantly underestimated in the atmospheric models. Nevertheless, the hierarchy holds anyway.  The mixing-length parameter probably decreases along the AGB, but its variation should not exceed 15\\,\\% per magnitude. This was taken into account in the abovementioned results.  This study confirms the findings in Paper I, the discrimination between Miras and semiregulars is not pertinent\\,: Groups 1 and 2 not only have similar kinematics but also the same pulsation mode.\\\\  It has also been shown that both the linear and nonlinear models that were the basis of all previous studies of LPV pulsation are probably far from the real pulsational behaviour of these stars. While dynamical calculations including a modern equation of state predict a strong reduction of the fundamental nonlinear period with respect to the linear one, important, {\\it positive} systematic corrections have to be applied to the periods of our linear models (30--45\\,\\% for the fundamental mode and 8--13\\,\\% for the first  overtone). Improvements of the physics of the sub-photospheric envelope  (phase lagged convection, turbulent pressure, horizontal opacity  averaging...) appear insufficient to explain these two shifts altogether, so  that the actual fundamental period should always exceed the theoretical one  by at least 15\\,\\%. This led us to conclude that all existing linear and  nonlinear pulsation codes probably suffer from neglecting  the stellar wind generated by the interaction with the circumstellar  envelope.\\\\  \\begin{figure*}[!t] \\vspace{7cm} \\special{psfile=MS9605f11a.eps voffset=-20 hoffset=-20 hscale=50 vscale=50} \\special{psfile=MS9605f11b.eps voffset=-20 hoffset=+250 hscale=50 vscale=50} \\caption[]{The PL ({\\bf a}) and PC ({\\bf b}) distributions of extended-disk/halo LPVs, compared to calibrated models with $Z=0.01$.} \\label{fig11} \\end{figure*}  As a consequence, the works of Barth\\`es \\& Tuchman (1994) and Barth\\`es \\& Mattei (1997), who confronted LNA models with the Fourier components of the lightcurves of a few nearby Miras and concluded in favour of the first overtone, should now be reconsidered by taking into account the necessary period corrections and the variations of metallicity and, possibly, mixing-length parameter.  This study may also have consequences in AGB evolutionary calculations, especially those that use the period as a substitute for a fundamental parameter ($M$ or $T_{\\rm eff}$) or as the variable in the empirical mass-loss function (e.g. Whitelock 1986, Vassiliadis \\& Wood 1993, Reid et al. 1995, Marigo et al. 1996). Indeed, these studies are based on pulsation models by Wood that appear to strongly overestimate the periods and their dependence on luminosity and metallicity, probably because of the equation of state. Concerning the fundamental mode, this peculiarity allows Wood's models to roughly mimic the abovementioned systematic period shifts and variation of the mixing length parameter, but with an uncertainty that remains to be assessed. Moreover, while many Long Period Variables are obviously pulsating on the first or second overtone, this possibility is usually neglected in evolutionary calculations, and Wood's models appear definitely inapropriate for these modes. These issues will be investigated in a forthcoming paper.\\\\  \\bigskip"
    },
    "0011/astro-ph0011208_arXiv.txt": {
        "abstract": "The nearby Seyfert galaxy NGC 7213 has been imaged in \\ha and HI with the CTIO 1.5 m telescope and with the Australia Telescope Compact Array (ATCA), respectively. Optically NGC 7213 looks undisturbed and relatively featureless but the continuum-subtracted \\ha image shows a 19 kpc long filament located approximately 18.6 kpc from the nucleus. The \\ha filament could be neutral gas photo-ionized by the active nucleus, as has been suggested for the Seyfert galaxy NGC 5252, or shock-ionized by a jet interacting with the surrounding HI, as has been suggested for the radio galaxy PKS 2240-41. The HI map reveals NGC 7213 to be a highly disturbed system suggesting a past merging event. ",
        "introduction": "NGC 7213 is a face-on Sa (Tully 1988) galaxy located at a distance of 22.0 Mpc ($H_{0}=75 kms^{-1}Mpc^{-1}$) (\\markcite{Tully1988}Tully 1988) and hosts an active nucleus.  Its nuclear activity was first discovered in follow-up optical spectroscopy of X-ray sources observed by the HEAO-A2  satellite (\\markcite{Phillips1979}Phillips 1978) and is in the \\markcite{Piccinotti1982}Piccinotti \\etal (1982) sample of X-ray selected active galaxies.  The nucleus of NGC 7213 has been classified as a Seyfert 1 (\\markcite{Phillips1979}Phillips 1979, \\markcite{FH1984}Filippenko and Halpern 1984) based on the broad optical emission lines, strong blue continuum and emission lines from such highly ionized species as Ne$^{+4}$. Its nuclear spectrum also exhibits narrow low ionization emission lines that define the LINER class which may originate from diffuse ionized gas surrounding the nucleus. This extended nuclear emission line region (ENER) has been observed in several spirals (e.g. M81,M31, N1398 etc.) and its ionization may be attributed to shock heating or photo-ionization by the bulge post asymptotic giant branch (PAGB) stars (\\markcite{Heckman1996}Heckman 1996; \\markcite{Devereux1997}Devereux, Ford, and Jacoby 1997).  The optical continuum image of the galaxy is dominated by an almost featureless bulge and shows no sign of any recent interaction (see Figure 5 of \\markcite{HD1999}Hameed \\& Devereux 1999).  An unsharp masking of the continuum image shows dust lanes sinuously winding into the central region which are also visible in the photograph of NGC 7213 in the {\\it Carnegie Atlas of Galaxies} (\\markcite{SB1994}Sandage and Bedke 1994). The continuum-subtracted \\ha image shows a ring of \\hii regions surrounding the nucleus that has been studied in detail by \\markcite{Storchi1996}Storchi-Bergmann \\etal (1996). However, the present paper is motivated by the discovery of a giant \\ha filament located approximately 18.6 kpc (in projection) south of the nucleus. This filament is roughly 19 kpc long, has no counterpart in the optical continuum image and lies well outside the optical diameter of the galaxy (\\markcite{HD1999}Hameed \\& Devereux 1999).  An HI map of NGC 7213, obtained as part of an HI survey of nearby Seyfert galaxies (\\markcite{Blank1999}Blank  1999), reveals NGC 7213 to be a highly disturbed system with tidal tails. The \\ha filament is, in fact, a small part of an HI tail located south of the nucleus. Contrary to the optical continuum image, the disturbed HI morphology suggests that NGC 7213 has gone through a merger or an interaction. In this paper we present details of neutral and ionized hydrogen gas morphology of NGC 7213. ",
        "conclusions": "{\\it What is the source of ionization for the giant filament?} It is unlikely that the filament is foreground galactic emission since our \\ha image was obtained with a redshifted \\ha filter and any foreground emission would have been subtracted out. In addition,  the \\ha morphology and its correlation with HI strongly suggests a connection with NGC 7213. The lack of clumpy HII regions in the filament also rule out local star formation as the ionization source. Similarly, there is no continuum emission from the filament, ruling out field OB and post-asymptotic Giant Branch stars as a source of ionization.  There are at least two possible mechanisms which might explain the large \\ha filament. One possibility is a starburst-related superwind. \\markcite{DB1999}Devine \\& Bally (1999) found a 3 kpc long ionized ``cap'' located 11 kpc above the plane of M82. Fainter \\ha was also detected between the \\ha cap and M82. Devine and Bally have suggested that the \\ha in the cap traces ambient material that is being shocked by a superwind and/or photoionized by radiation from nuclear starburst. Massive star formation in NGC 7213 is confined to a ring of HII regions surrounding the nucleus and is much weaker than the starburst in M82. Our calculations suggest that it is very unlikely that the filament in NGC 7213, which is much larger than the ``cap'' in M82(Table 1), is being ionized by the leakage of UV photons from the star forming ring.  A more likely explanation is that the filament is either being photo-ionized by UV photons escaping from the central AGN or shock-ionized by a jet interacting with the surrounding medium. Observations of Seyfert galaxies have identified galactic scale minor outflows in a number of galaxies (e.g. \\markcite{WT1994}Wilson \\& Tsvetanov 1994; \\markcite{Colbert1996}Colbert \\etal. 1996). The presence of broad emission lines (\\markcite{FH1984}Filippenko \\& Halpern 1984), along with the detection of X-rays(\\markcite{Piccinotti1982}Piccinotti et al. 1982; \\markcite{Fabbiano1992}Fabbiano, Kim, and Trinchieri 1992) and an unresolved UV point source (\\markcite{RB1999}Rokaki \\& Boisson 1999) in the nucleus of NGC 7213, strongly suggest  the presence of an AGN. Lyman continuum photons from the central engine may ionize part of the extensive HI which surrounds NGC 7213. Such a situation exists within the type 1 Seyfert galaxy NGC 5252, which hosts a biconical structure of ionized gas that extends up to a distance of 18 kpc from the nucleus and consists of a complex network of filaments(\\markcite{PF1996}Prieto \\& Freudling 1996; \\markcite{Tsvetanov1996}Tsvetanov \\etal 1996; \\markcite{Morse1998}Morse \\etal 1998). The largest of these filaments is located about 7 kpc from the nucleus and is approximately 3.6 kpc in length (Table 1).  In the case of NGC 7213, the \\ha luminosity of the nucleus is $4.7 \\times 10^{40} \\lum$. If the filament in NGC 7213 is photoionized by the nucleus, then, at least, 14$\\%$ of ionizing photons have to escape from the nucleus to provide the observed \\ha luminosity of the filament: $6.4 \\times 10^{39} \\lum$. We can obtain a rough estimate as to whether or not the source responsible for \\ha emission in the nucleus has enough energy to ionize the filament. Lets consider a simple case where ionizing photons are leaking isotropically from the nucleus. The flux at the distance of of the filament (18.6 kpc), $F_{nf}$, can be calculated by $F_{nf}=L(H\\alpha)_{nuc}/4\\pi D^{2}$, where $L(H\\alpha)_{nuc}$ is \\ha luminosity of the nucleus  and D is the distance of the filament from the nucleus. Using the above values, we get $F_{nf}=1.1 \\times 10^{-6} \\flux$. Assuming the filament to be a 2-dimensional sheet, we can get the luminosity of the filament, $L_{nf}$, by multiplying $F_{nf}$ by the surface area of the filament, $165 \\arcsec \\times 55 \\arcsec$ or $5.9 \\times 10^{22}cm \\times 1.8 \\times 10^{22}cm$. This gives $L_{nf} = 1.2 \\times 10^{39} \\lum$ which is about 6 times smaller than the observed luminosity of the filament (Table 1). It should be noted, however, that the \\ha luminosity of the nucleus is not an accurate measure of the ionizing flux, and represents only those ionizing photons that have been absorbed by hydrogen near the nucleus. The covering factor for the broad line region in AGNs is quite uncertain, but some studies (e.g. \\markcite{Peterson1997}Peterson 1997) estimate the covering factor to be around 0.1, suggesting that only 10$\\%$ of the ionizing photons are being absorbed. If this is the case in NGC 7213, then the nucleus can easily provide the photons required to ionize the filament. In addition, our calculations above used the simplest case of isotropic emission. Even a modest beaming of the ionizing photons towards the filament can potentially provide enough energy to ionize the filament.  On the other hand, the radio galaxy PKS 2250-41 provides an example where a giant emission line arc is shock ionized by a radio jet(\\markcite{Tadhunter1994}Tadhunter \\etal 1994). It should be noted, however, that the emission line arc in PKS 2250-41 has been observed only in [OIII] and no \\ha observations for the galaxy exist in the literature. The arc in PKS 2250-41 is located 37 kpc from the nucleus and has a total length of about 50 kpc (\\markcite{Clark1997}Clark \\etal 1997). Such jet interaction induced ionization could be the case in NGC 7213, albeit on a smaller scale. In the study of star forming regions of NGC 7213, \\markcite{Storchi1996}Storchi-Bergmann \\etal (1996) found velocity dispersions near the nucleus suggesting a collimated outflow or other non-circular motions. Furthermore, \\markcite{Harnett1987}Harnett (1987) and \\markcite{Blank1999}Blank (1999), found extended radio emission in low resolution images that might also indicate the presence of a radio jet.  Understanding the kinematics of the ionized gas could help distinguish between these two possibilities. If ionized gas is produced by UV photons escaping from the nucleus, then we might expect the ionized gas to have the same radial velocity as the neutral gas at the same location. On the other hand, if the ionized gas is produced by shocks, its velocity should be significantly different from the neutral gas. In addition, optical spectroscopy can also provide diagnostic line ratios that may determine whether the filament is shock-ionized or photo-ionized. Shock-ionization is expected to produce high ratios of [NII](6548+6584$\\AA$)/\\ha,[SII](6716+6731$\\AA$)/\\ha, and [OIII](5007$\\AA$)/\\ha relative to photo-ionization (\\markcite{SM1979}Shull \\& McKee 1979; \\markcite{DS1995}Dopita \\& Sutherland 1995).  The identity of the progenitor galaxy, responsible for the disturbed HI morphology of NGC 7213, is also unclear.  There is a nearby elliptical galaxy of unknown redshift located at $5\\arcmin$ north-east of NGC 7213. If that system is a close companion, it is unlikely to be responsible for the disturbed HI morphology. While the anomalous HI arms  could be produced by a small galaxy tidally removing HI gas from its larger companion (e.g. the M 51 system), the elliptical companion of NGC 7213 is located in the middle of the northern arm, not at the end of the arm as one would expect, if it was drawing out the HI. It is more likely that NGC 7213 is a merger remnant, as \\markcite{Blank1999}Blank (1999) has identified an inner HI disk counter-rotating with respect to the larger HI disk in NGC 7213. This notion is further supported by the presence of ``shells'' in an unsharp mask of the Digital Sky Survey image of NGC 7213 (Hibbard, private communication). However, we cannot rule out the presence of another perturber, as there is a bright foreground star to the north-west of NGC 7213 that may be masking any nearby galaxy.  The giant filament in NGC 7213, one of the largest among nearby active galaxies, joins an increasing number of galaxies that have galactic scale ionization features. It is difficult to estimate how common these objects are, as very few studies of Seyfert galaxies exist that have focused on large scale features. In a study of 22 edge-on Seyfert galaxies, \\markcite{Colbert1996}Colbert et al. (1996) find $\\sim 1/4$ of galaxies with good evidence for minor axis outflows, but none at scales comparable to NGC 7213 and NGC 5252. Similarly, \\markcite{Baum1993}Baum et al. (1993) found kpc-scale extended features in seven out of ten Seyfert galaxies they studied. A large, wide field, systematic study of nearby Seyfert galaxies will provide information about the importance of these filaments and their relationship to AGNs and their host galaxies."
    },
    "0011/astro-ph0011178_arXiv.txt": {
        "abstract": "We present spatially resolved near-infrared and mid-infrared (1--25~\\microns) imaging of the WL~20 triple system in the nearby (\\( d=125 \\)~pc) \\( \\rho \\) Ophiuchi star-forming cloud core. We find WL~20 to be a new addition to the rare class of {}``infrared companion systems{}'', with WL~20:E and WL~20:W displaying Class II (T-Tauri star) spectral energy distributions (SEDs) and total luminosities of 0.61 and 0.39~\\( L_{\\sun } \\), respectively, and WL~20:S, the infrared companion, with a Class I (embedded protostellar) SED and a luminosity of 1.0--1.8~\\( L_{\\sun } \\). WL~20:S is found to be highly variable over timescales of years, to be extended (40~AU diameter) at mid-infrared wavelengths, and to be the source of the centimeter emission in the system.  The photospheric luminosities of 0.53~\\( L_{\\sun } \\) for WL~20:E and 0.35~\\( L_{\\sun } \\) for WL~20:W, estimated from our data, combined with existing, spatially resolved near-infrared spectroscopy, allow us to compare and test current pre-main-sequence evolutionary tracks. The most plausible, non-accreting tracks describing this system are those of \\citet{dantona}. These tracks give an age of 2--\\( 2.5\\times 10^{6} \\)~yr and masses of 0.62--0.68~\\( M_{\\sun } \\) for WL~20:E and 0.51--0.55~\\( M_{\\sun } \\) for WL~20:W, respectively. The age and mass of WL~20:S cannot be well determined from the currently available data. WL~20:E and WL~20:W fall into the region of the H-R diagram in which sources may appear up to twice as old as they actually are using non-accreting tracks, a fact which may reconcile the co-existence of two T-Tauri stars with an embedded protostar in a triple system. The derived masses and observed projected separations of the components of the WL~20 triple system indicate that it is in an unstable dynamical configuration, and may therefore provide an example of dynamical evolution during the pre-main-sequence phase. ",
        "introduction": "The optically undetectable source, WL~20 \\citep[also known as BKLT J162715\\(-\\)243843 and GY~240, see][for other aliases]{bklt}, was discovered in a near-infrared (near-IR) bolometer survey of a \\( 10\\arcmin \\times 10\\arcmin \\) region of self-absorbed \\( ^{13} \\)CO emission in the \\( \\rho \\) Ophiuchi star-forming cloud \\citep{wl}. Soon thereafter, WL~20 was detected at 10~\\microns\\ from ground-based observations \\citep{lw}, and at longer wavelengths in the pointed observations mode of the \\textit{IRAS} satellite, where it is referred to as YLW~11 \\citep{ylw}.  In the currently accepted classification scheme of young stellar objects (YSOs) devised by \\citet{lada}, WL~20 was one of the first sources to be identified as a Class I source \\citep{wly}. Empirically, Class I sources have broader than blackbody spectral energy distributions (SEDs), with rising 2--10~\\micron\\ spectral slopes, \\( a>0.3 \\) (where \\( a=\\frac{d\\log \\lambda F_{\\lambda }}{d\\log \\lambda } \\)). Theoretically, the SEDs of Class I sources are interpreted to correspond to a remnant infalling dust and gas envelope surrounding a central protostar\\( + \\)disk system \\citep{als}. The more evolved Class II sources, with \\( -0.3>a>-1.6 \\), are pre-main-sequence (PMS) star\\( + \\)disk systems that have dispersed their remnant infall envelopes.  With the advent of near-IR array detectors, WL~20 was soon resolved into a binary system at 2.2~\\microns\\ using a pixel scale of 0\\farcs85 \\citep{RAB} and first reported to have an east-west separation of \\( \\approx \\) 2\\farcs7 from observations with a pixel scale of 0\\farcs78 \\citep{Bar89}. This separation corresponds to \\( \\approx \\) 340~AU for an adopted distance to the cloud of 125~pc.  We note that some confusion still exists as to the distance to the \\( \\rho \\) Ophiuchi clouds, primarily due to many authors citing the distance to the adjacent Sco-Cen OB association instead. The distance to the Sco-Cen OB association has been determined to be 160\\( \\pm \\)10~pc photometrically \\citep{whit} or 145~pc astrometrically \\citep{deZ99}. By contrast, the distance to the \\( \\rho \\) Ophiuchi cloud itself has been determined to be 125\\( \\pm \\)25~pc from detailed kinematic studies of the cloud gas \\citep{deG89,deG92}, and, more recently, by a study of \\textit{Hipparcos} parallaxes and Tycho \\( B-V \\) colors of stars of Classes III \\& V, which show an abrupt rise in reddening at \\( d=120 \\)~pc, as expected for a molecular cloud \\citep{knu}. We therefore adopt a distance of 125~pc to the cloud in this paper.  The first indication that WL~20 is a triple system came from a deep ProtoCAM survey to identify near-IR counterparts of centimeter continuum sources in the \\( \\rho  \\) Ophiuchi cloud core \\citep{sks}. The VLA source identified with WL~20 is known as LFAM~30, and was imaged at 6~cm with an 11\\arcsec \\( \\times  \\)5\\arcsec\\ beam \\citep{LFAM}. With the sensitive ProtoCAM images acquired at a pixel scale of 0\\farcs20, a third and weakest 2.2~\\micron\\ component of the system was easily identified, and designated as {}``30S{}'' (for the Southern component of LFAM~30), whereas the components of the previously known near-IR binary are referred to as {}``30E{}'' and {}``30W{}'' by these authors \\citep{sks}. With 30E as the positional reference, 30W is quoted at a separation of 3\\farcs3 at P.A. 269\\degr, and 30S at a separation of 3\\farcs9 at P.A. 232\\degr.  Near-IR spectra of the two brighter \\( K \\) components of the WL~20 triple system, WL~20:E (\\( K=10.13 \\)) and WL~20:W (\\( K=10.40 \\)) were presented as part of a spectroscopic survey of YSOs with \\( K<10.5 \\) in \\( \\rho \\) Oph \\citep{gl}. The near-IR spectra of both WL~20:E \\& W were found to be consistent with those of other Class II sources, with a K--M spectral type established for WL~20:E, and a more precise K7--M0 spectral type determination for WL~20:W, due to the less severe continuum veiling in its spectrum \\citep{gm}. At their assumed distance of 160~pc, WL~20:E was determined to have a bolometric luminosity, \\( L_{bol}=0.4L_{\\sun } \\), through \\( A_{V}=15.4 \\), whereas WL~20:W was determined to have \\( L_{bol}=1.7L_{\\sun } \\) through \\( A_{V}=18.1 \\). Continuum veiling in the spectrum of WL~20:E precluded its placement on the H-R diagram, whereas the lesser veiling in the spectrum of WL~20:W allowed a mass estimate of 0.3~\\( M_{\\sun } \\) from its location along a \\( 3\\times 10^{5} \\)~yr pre-main-sequence isochrone \\citep{gm}.  A more recent spectroscopic survey of near-IR sources in the \\( \\rho \\) Ophiuchi core including many sources as faint as \\( K\\sim 12 \\), was obtained at higher spectral resolution (\\( R=1200 \\)) than the previously published surveys (with \\( R\\leq 1000 \\)), allowing for significantly improved classifications for G through M spectral types \\citep{lr}. These authors assigned a K6 spectral type to WL~20:E (GY 240B), with a bolometric luminosity of 0.55~\\( L_{\\sun } \\), and an M0 spectral type to WL~20:W (GY 240A) with \\( L_{bol}=1.4L_{\\sun } \\) (they also assumed a distance of 160~pc). Both sources were found to have a foreground extinction of \\( A_{V}=16.3 \\). WL~20:S, with \\( K\\sim 12.6 \\), was excluded from this survey as well, due to its relative dimness at near-IR wavelengths.  Interestingly, the discrepancy between the Class I SED classification of WL~20 on the one hand, and its near-IR Class II spectroscopic classifications on the other, has not been remarked upon previously. This inconsistency can only be addressed by producing spatially resolved SEDs of the individual components of this triple system. Until now, the highest spatial resolution photometry of WL~20 longward of 4.8~\\microns\\ has been through a 6--8\\arcsec\\ aperture---confusing all three components \\citep{lw}. In order to better constrain the properties of WL~20:S, and of the triple system of which it is a member, we have obtained new, unprecedentedly high (sub-arcsecond) spatial resolution, ground-based mid-IR images of the WL~20 system, at six separate wavelengths spanning the 8--25~\\micron\\ atmospheric window. Additionally, we present spatially resolved near-IR imaging of this triple system. Finally, we have performed careful astrometry, allowing us identify the source of the centimeter continuum emission. ",
        "conclusions": "We have presented sub-arcsecond, mid-infrared imaging photometry of the WL~20 triple system at 7.9, 10.3, 12.5, 17.9, 20.8, and 24.5 \\( \\mu \\)m. When supplemented by spatially-resolved, near-infrared imaging photometry from ProtoCAM at the IRTF, these combined data allow solid determinations of the spectral energy distribution of each source individually, as well as accurate luminosity determinations. We find the source luminosities for WL~20:E, WL~20:W, and WL~20:S to be 0.61~\\( L_{\\sun } \\), 0.39~\\( L_{\\sun } \\), and 1.0--1.8~\\( L_{\\sun } \\), respectively. For WL~20:E and WL~20:W, 0.53~\\( L_{\\sun } \\) and 0.35~\\( L_{\\sun } \\) can be attributed to photospheric emission alone.  WL~20 can now be classified as an {}``infrared companion system,{}'' with WL~20:S exhibiting an embedded protostellar (Class I) SED, while its two neighbors, WL~20:E and WL~20:W, each exhibit T~Tauri star (Class II) SEDs. The infrared companion, WL~20:S, is the dominant luminosity source in the system. WL~20:S differs from its T~Tauri companions in three important respects: 1) its near- and mid-IR fluxes vary significantly over timescales of years; 2) it is well-resolved at mid-IR wavelengths, with a constant, wavelength-independent source diameter of 40~AU; and 3) it is found to be the source of the radio cm emission in the system.  Since the effective temperatures of WL~20:E \\& W are known from spatially-resolved near-IR spectroscopy, we can place these sources on a Hertzsprung-Russell diagram. We are thus able to test currently available pre-main-sequence evolutionary tracks at unprecedentedly high temporal resolution and find that of the non-accreting models, the \\citet{dantona} tracks yield the most plausible system age, at 2.0--\\( 2.5\\times 10^{6} \\)~yr. The inferred source masses from these tracks at these ages are 0.62--0.68~\\( M_{\\sun } \\) for WL~20:E and 0.51--0.55~\\( M_{\\sun } \\) for WL~20:W, respectively.  We cannot, at present, independently determine a mass or age for the infrared companion, WL~20:S. However, the intriguing possibility now exists of determining the spectral type, and, therefore, the effective temperature of this embedded source with the new generation of high-resolution spectrographs on 8--10~m ground-based telescopes. Once an effective temperature determination has been made spectroscopically, and assuming system coevality, one could locate WL~20:S on an isomass track, independently of its known luminosity. Thus, the possibility exists, for the first time, to directly derive the accretion luminosity of a Class I protostar.  Millimeter interferometry of this unique triple system would advance our understanding of the gas dynamics involved, processes which cannot be explored in any other way. Higher temporal resolution spatially-resolved imaging and monitoring of WL~20 at infrared wavelengths, combined with detailed modeling of its appearance will also lead to a more detailed understanding of the actual accretion processes taking place in WL~20:S."
    },
    "0011/astro-ph0011452_arXiv.txt": {
        "abstract": "We study Sunyaev-Zel'dovich Effect (SZE) cluster counts in different cosmologies. It is found that even without the full knowledge of the redshift distribution of SZE clusters, one can still readily distinguish a flat universe with a cosmological constant from an open universe. We divide clusters into a low redshift group (with redshift $z\\le 0.5$) and a high redshift group (with $z\\ge 1$), and compute the ratio of $r=N(z\\le 0.5) /N(z\\ge 1)$, where $N(z\\le 0.5)$ is the number of flux-limited ($S_{\\nu}^{lim}$) SZE clusters with $z\\le 0.5$ and $N(z\\ge 1)$ is the number of flux-limited SZE clusters with $z\\ge 1$. With about the same total number of SZE clusters $N(z\\ge 0)$, the $r$ value for a flat universe with a non-zero cosmological constant and that for an open universe occupy different regions in the $S_{\\nu}^{lim}-r$ plot for the most likely cosmological parameters $0.25\\le \\Omega_0\\le 0.35$ and $0.2\\le \\Gamma\\le 0.3$, where $\\Omega_0$ is the matter density parameter of the universe, and $\\Gamma$ is the shape parameter of the power spectrum of linear density fluctuations. Thus with a deep SZE cluster survey, the ratio $r$ can reveal, independent of the normalization of the power spectrum, whether we are living in a low-density flat universe or in an open universe. Within the flat universe scenario, the SZE cluster-normalized $\\sigma_8$ is studied, where $\\sigma_8$ is the r.m.s. density fluctuation within the top-hat scale $8\\hbox { Mpc}h^{-1}$ where $h$ is the Hubble constant in units of $100\\hbox { kms}^{-1}\\hbox { Mpc}^{-1}$. A functional relation $\\sigma_8\\propto \\Omega_0^{-0.13}$ is found. Combined with the X-ray cluster-normalized $\\sigma_8\\propto \\Omega_0^{-0.52+0.13\\Omega_0}$, one can put constraints on both $\\Omega_0$ and $\\sigma_8$ simultaneously. ",
        "introduction": "Clusters of galaxies are the largest virialized objects in the universe, and contain valuable information of the universe and of the large-scale structure. There have been intensive studies on clusters from different approaches, such as strong and weak gravitational lensing, X-ray, and optical observations. With the technical advents of interferometers, the cluster's Sunyaev-Zel'dovich Effect (SZE) (Sunyaev \\& Zel'dovich 1970, 1980; Birkinshaw 1999, Carlstrom et al. 1999), a spectral distorsion of Cosmic Microwave Background Radiation (CMB) due to scattering of CMB photons by hot electrons within clusters, has been becoming a new probe for the cluster study. As CMB photons pass through intracluster hot electrons, on average they gain energies through the inverse-Compton scattering, and as a result of this, the number of low energy photons decreases while the number of high energy ones increases. Thus for observations with the frequency $\\nu > (<) 219\\hbox { GHz}$, hot clusters behave like emitting sources (absorbers) of photons. The equivalent temperature increment (or decrement) $\\Delta T$ of CMB photons toward clusters is proportional to $\\int n_e T_{gas} dl$, where $n_e$ is the number density of electrons, $T_{gas}$ is the hot gas temperature, and $dl$ is the line element along the line of sight. The integrated SZ effect of a cluster is then directly proportional to the cluster's gas mass if the gas is close to be isothermal. Therefore the integrated SZ effect is not sensitive to the lumpy structures of the gas and the gas fraction can be estimated relatively clean from the SZ effect in conjunction with lensing observations (e.g., Grego et al. 2000). On the other hand, because of the different dependence on $n_e$ of the cluster's X-ray surface brightness ($S_x\\propto \\int n_e^2 \\Lambda_{eH} dl$, where $\\Lambda_{eH}$ is the X-ray cooling function) and of the SZ effect, the angular diameter distance to a cluster can be derived directly from a joint analysis of the X-ray emission and the SZ effect through modeling the gas density profile properly (e.g., Reese et al. 2000).  Apart from the SZE studies for individual clusters, statistical investigations on large number of SZE clusters can also yield very promising results. In fact several interferometric arrays have been proposed for surveying SZE clusters, including the AMIBA (Array for MIcrowave BAckground) project which has been founded in Taiwan. Due to the frequency (energy) redshift dependence of the CMB photons, the SZ effect is independent of the redshift, which permits a SZE cluster survey to detect very high redshift clusters with relative ease. Therefore the cluster redshift evolution can be studied with high statistical significance. Another advantage of the SZE cluster studies over those of X-ray is that the integrated SZ effect of an individual cluster is directly proportional to the gas mass within the cluster (assuming the gas is isothermal), which is in turn related to the total mass of the cluster, and the number of SZE clusters can be predicted analytically from the Press-Schechter formula (or other similar models) in a straightforward manner with certain qualifications. By contrast, because of the $n_e^2$ dependence of the X-ray surface brightness, the gas density profile has to be modeled in the X-ray studies to estimate flux-limited X-ray cluster number counts analytically, a procedure that can introduce large uncertainties (the prediction on the number of temperature-limited X-ray clusters suffers less problem). One may investigate both aspects of clusters with numerical simulations, but finite numerical resolutions and box sizes ultimately limit their applications, and the analytical analysis can be complementary.  Among other promising aspects, the redshift distribution of SZE clusters can be used to constrain cosmological parameters. The redshift distribution of SZE clusters with $\\Omega_0=1$ is distinctly different from those with $\\Omega_0\\sim 0.3$, and the existence of several SZE clusters at redshift $z\\gsim 1$ would strongly exclude the $\\Omega_0=1$ model. The difference between the redshift distribution of a low-density flat universe model and of a low-density open universe model is less dramatic, and one needs a relatively large number of clusters at high redshift to falsify them. The SZE cluster surveys are suitable for this purpose. By fully using the redshift distribution and the total number of SZE clusters, Haiman, Mohr, \\& Holder (2000) studied constraints on the quintessence theory from future SZE (and X-ray) cluster surveys. The redshift of a cluster with $z\\le 1$ can be obtained at least by using the photometric method around a characteristic spectral break. For clusters with $z>1$, it however appears difficult to measure the redshifts precisely except for very large clusters. In this paper, we propose a method which can distinguish, within the parameter regime $0.25\\le \\Omega_0\\le 0.35$ and $0.2\\le \\Gamma \\le 0.3$, a flat universe with a non-zero cosmological constant from an open universe even without knowing the full redshift distribution of SZE clusters. We divide clusters into two groups: a low redshift one with $z\\le 0.5$, and a high redshift one with $z\\ge 1$. We study the ratio $r=N(z\\le 0.5)/N(z\\ge 1)$ for different cosmologies with different parameters, where $N(z\\le 0.5)$ is the total flux-limited number of clusters with $z\\le 0.5$, and $N(z\\ge 1)$ is the total flux-limited number of clusters with $z\\ge 1$. It is found that $r$ can be used to disentangle a low-density flat universe from a low-density open universe. Notice that to compute $r$, we only need to know the redshift range of a cluster (whether it is $z\\le 0.5$ or $z\\ge 1$) rather than its precise redshift.  We also investigate, within the framework of the flat universes with a non-zero cosmological constant, the SZE cluster-normalized $\\sigma_8$. The $\\sigma_8-\\Omega_0$ relation inferred from the SZE cluster counts is distinctly different from that from the X-ray cluster counts with the latter more steeper. This difference can be used to limit the parameter regime of $\\Omega_0$ and $\\sigma_8$ simultaneously.  The rest of the paper is organized as follows. Section 2 will present the formulation for the study. The results will be shown in section 3. Section 4 contains a summary. ",
        "conclusions": "In our analyses, we used $f_{ICM}=0.1$, $X=0.76$ and $[d \\hbox {ln} \\rho_{gas}(r)/ d\\hbox {ln} r]_{r_{vir}}=2$. To change these parameters however, is equivalent to change the overall flux limit $S_{\\nu}^{lim}$. From the figures we showed, it is easy to see that all our conclusions remain qualitatively unchanged for different choices of these parameters. We have assumed the hydrodynamic equilibrium and isothermality for the intracluster gas. The Press-Schechter formalism has been adopted to calculate the cluster counts.  We studied the $r$ quantity for two types of cosmologies: the low-density flat models and the low-density open models. Within the studied parameter regimes $0.25\\le \\Omega_0 \\le 0.35$ and $0.2\\le \\Gamma \\le 0.3$, the $r$ value for the two sets of cosmologies are well separated. The flat model with $\\Omega_0=0.3$ and $\\Gamma=0.25$ and the open model with the same $\\Omega_0$ and $\\Gamma$ can be differentiated at the $3\\sigma$ level for a $50 \\hbox { deg}^2$ survey. Since we normalize different models in a way such that they give rise to the same total number of SZE clusters, our analyses naturally take into account the total number of SZE clusters. For wider parameter ranges, the information from other observations, such as X-ray or optical cluster surveys has to be used to differentiate the two types of models.  There are other ways to determine the geometry of the universe. Measurements on fluctuations of the CMB radiation provide a clean test on this aspect (e.g., Hu 2000). The horizon size of the universe at decoupling separates large-scale and small-scale CMB fluctuations. Large-scale fluctuations were outside the horizon when photons escaped while small-scale perturbations were within the horizon at decoupling and therefore sustained acoustic oscillations. The position of the primary Doppler peak of the CMB fluctuation power spectrum is determined by the angular size of the horizon at decoupling. For a flat universe (low-density with a non-zero cosmological constant or high density), the peak located at $l\\approx 200$, and the peak is shifted to smaller angular scale or higher $l$ for a low-density open universe, where $l$ represents the two-dimensional angular wavenumber. Observations have seen the rising and the declining of CMB fluctuations around $l\\sim 200$ (e.g., Miller et al. 1999, de Bernardis et al. 2000), which constrains convincingly that $\\Omega_{tot}$ is close to $1$, where $\\Omega_{tot}$ is the total density parameter. Combining with other information, e.g., supernova measurements (Perlmutter et al. 1999, Schmidt et al. 1998), leads to the conclusion that we are living in a low-density flat universe. If the results from SZE cluster surveys are in agreement with the CMB results, we will be in a more solid position to say that the universe is flat and the simple structure formation theory we adopted here is reasonably correct. On the other hand, any inconsistency between the SZE cluster results and the results from CMB measurements would pose challenges to our understanding of the universe. For example, alternate structure formation theories, such as non-Gaussian initial fluctuation models, may need to be considered.  Should the universe be flat, SZE cluster surveys can also provide constraints on cosmological parameters. We studied $\\sigma_8(\\Omega_0)$ inferred from the total SZE cluster counts. A functional relation $\\sigma_8\\propto\\Omega^{-0.13}$ is found. Combined with the current X-ray cluster-normalized $\\sigma_8$, the parameters can be determined (take $\\Omega_0=0.3$ and $\\sigma_8=0.93$ as the central values) to $\\Omega_0=(1\\pm 27\\%) \\times 0.3$ and $\\sigma_8=(1\\pm 3.2\\%)\\times 0.93$ for the SZE cluster counts confined to the $3\\sigma$ level in a $50\\hbox { deg}^2$ survey. Note that these constrains are from the total number of clusters only, and no redshift information is needed."
    },
    "0011/astro-ph0011387_arXiv.txt": {
        "abstract": "Dusty galaxies with luminosities in excess of $10^{11}$\\,L$_\\odot$ have been detected out to redshifts $z \\sim 1$ by the {\\it Infrared Space Observatory (ISO)}, and to higher redshifts using millimetre(mm)- and submm-wave cameras on ground-based telescopes. The integrated properties of these more distant galaxies are also constrained by measurements of the intensity of the submm-wave background radiation. While it is generally unclear whether their energy is released by gravitational accretion or by star formation, circumstantial evidence favours star formation. Unless these high-redshift galaxies are extremely massive, which is not expected from standard models of galaxy evolution, this luminosity cannot be sustained for more than a fraction of a Hubble time, and so they are undergoing some sort of `bursting' behaviour. The interpretation and analysis of this population is discussed, and the key observations for deriving a robust history of their evolution, which is likely to be the history of starburst activity, are highlighted. ",
        "introduction": "From observations of low-redshift dusty galaxies using the {\\it IRAS} satellite \\cite{SandN} close to the peak of their restframe spectral energy distributions (SEDs), it is known that a similar amount of energy in the local Universe is produced by stars in dust-enshrouded and dust-free environments. The comoving luminosity density of dusty galaxies is also known to evolve strongly, from the slope of the faint counts of {\\it IRAS} galaxies at 60\\,$\\umu$m \\cite{Bertin}, which provide information to $z \\simeq 0.2$. {\\it ISO} observations at both shorter \\cite{Elbaz} and longer \\cite{Juvela} wavelengths confirm that strong evolution continues to $z \\sim 1$.  At longer wavelengths, the redshifted emission from very luminous, high-redshift dusty galaxies can be detected in the mm and submm wavebands. Independent surveys made using the 450/850-$\\umu$m SCUBA camera at the JCMT \\cite{Smail_catalogue} have determined the counts of high-redshift dusty galaxies. 1.2-mm surveys using the MAMBO detector array at the IRAM 30-m telescope \\cite{Carilli} have detected a similar population of galaxies. In three cases, the detection of CO emission from gas located at the position and redshift of a suspected optical identification (at $z=1.06$, 2.55 and 2.80) \\cite{Frayer} provides absolute confirmation of the identification. Extremely deep VLA radio images of the survey fields can be used to impose constraints on the redshifts and SEDs of the detected galaxies \\cite{CandY,SIOBK}. It is likely that the detected galaxies are at $\\bar z \\simeq 2-3$, and there are very few plausible low-redshift ($z \\le 1$) counterparts. The counts and redshift distributions of these distant dusty galaxies can be used to constrain models of galaxy evolution at high redshifts.  In addition to the detection of individual submm-selected galaxies, the background radiation intensity in the mm and submm wavebands \\cite{Fixsen} traces the integrated emission from the entire population of dusty high-redshift galaxies. The mm/submm-wave background spectrum has the form $I_\\nu \\propto \\nu^{2.64}$, and there is no clear spectral break down to 2000\\,$\\umu$m. The lack of a break to a steeper slope supports the idea that high-redshift galaxies with redshifted SEDs peaking at about 1000\\,$\\umu$m are still contributing to the background intensity, indicating a high maximum redshift of the population of about 10. If the shape of the SED is assumed not to evolve significantly, then the luminosity density must evolve as $\\rho_{\\rm L} \\propto (1+z)^{\\simeq -1.1}$ \\cite{BL93b,BJ} for $z \\gg 1 $ in order to generate this background spectrum, a result which is independent of cosmology. Such a gently declining high-redshift luminosity density is naturally consistent with the incomplete redshift distributions of galaxies detected in mm/submm-wave surveys.  Note that any model of the evolution of dusty galaxies must predict the redshift distribution of submm-detected galaxies correctly. Submm-wave surveys are very sensitive to high-redshift galaxies \\cite{BL93a}, and it is easy to propose models that fit both the observed submm-wave counts and background radiation spectrum at the expense of a redshift distribution that is biased far too high. The verification of predicted redshift distributions is therefore a crucial test of such models.  In this paper, the forms of evolution of galaxies that were previously derived in the context of far-infrared(IR) and submm-wave data are updated to take account of the much greater amount of information that has become available, particularly from {\\it ISO} surveys. The results are similar to, but less uncertain than, those derived earlier \\cite{BSIK,BJ}. The results have also been updated to include the currently favoured cosmological parameters $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ are assumed. ",
        "conclusions": "These models, which involve a minimum number of free parameters, provide a reasonable description of all the counts and redshift distributions of dusty galaxies at both high and low redshifts. As additional data, especially more complete redshift distributions and very deep {\\it SIRTF} mid-IR counts, become available, more details can be incorporated into the models to reveal further information about the properties of evolving distant dusty starbursts/AGN.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=.65\\textwidth,angle=-90]{blain_fig1.cps} \\end{center} \\caption[]{The history of star formation inferred using both methods discussed in the text. The thick solid and dashed lines represent the simple luminosity evolution model and the hierarchical model respectively. The thinner solid lines show the approximate envelope of 68\\% uncertainty in the results of the simple model. The data points are taken from a variety of sources: references can be found elsewhere \\cite{BJ,vdW}. The thin and thick dotted lines represent the best-fitting results obtained in previous derivations, the modified Gaussian \\cite{BSIK,Barger} and 35-K hierarchical models \\cite{BJ} respectively. The absolute normalisation of the curves depends on the assumed stellar initial mass function and the fraction of the dust-enshrouded luminosity of galaxies that is generated by AGN. } \\label{eps1} \\end{figure}  \\subsubsection*{Acknowledgements} The author, Raymond and Beverly Sackler Foundation Research Fellow at the IoA, thanks the Foundation for generous financial support, ESO for support at the meeting, and Vicki Barnard and Kate Quirk for helpful comments on the manuscript. These results are updated from SCUBA Lens Survey work with Ian Smail, Rob Ivison and Jean-Paul Kneib."
    },
    "0011/astro-ph0011102_arXiv.txt": {
        "abstract": "The lens mass determined from the photometrically obtained Einstein time scale suffers from large uncertainty due to the lens parameter degeneracy. The uncertainty can be substantially reduced if the mass is determined from the lens proper motion obtained from astrometric measurements of the source image centroid shifts, $\\deltavec\\thetavec_{\\rm c}$, by using high precision interferometers from space-based platform such as the {\\it Space Interferometry Mission} (SIM) and ground-based interferometers soon available on several 8 -- 10 m class telescopes.  However, for the complete resolution of the lens parameter degeneracy it is required to determine the lens parallax by measuring the parallax-induced deviations in the centroid shifts trajectory, $\\Deltavec\\deltavec\\thetavec_{\\rm c}$.  In this paper, we investigate the detectabilities of $\\deltavec \\thetavec_{\\rm c}$ and $\\Deltavec\\deltavec \\thetavec_{\\rm c}$ by determining the distributions of the maximum centroid shifts and the average maximum deviations expected for different types of Galactic microlensing events caused by various masses.  From this investigation, we find that as long as source stars are bright enough for astrometric observations it is expected that $\\deltavec \\thetavec_{\\rm c}$ for most events caused by lenses with masses greater than $0.1\\ M_\\odot$ regardless of the event types can be easily detected from observations by using not only the SIM (with a detection threshold $\\delta\\theta_{\\rm th}\\sim 3\\ \\mu{\\rm as}$) but also the ground-based interferometers (with $\\delta \\theta_{\\rm th}\\sim 30\\ \\mu{\\rm as}$).  However, detection of $\\Deltavec \\deltavec \\thetavec_{\\rm c}$ from ground-based observations will be difficult for nearly all Galactic bulge self-lensing events, and will be restricted only for small fractions of disk-bulge and halo-LMC events, for which the deviations are relatively large.  From observations by using the SIM, on the other hand, detecting $\\Deltavec \\deltavec\\thetavec_{\\rm c}$ will be possible for majority of disk-bulge and halo-LMC events and even for some fraction of bulge self-lensing events.  For the complete resolution of the lens parameter degeneracy, therefore, SIM observations (or equivalent) will be essential. ",
        "introduction": "The biggest difficulty in revealing the nature of Galactic dark matter by using microlensing is that one cannot determine the individual lens masses because the Einstein time scale, which is the only quantity related to the lens mass obtained from the photometrically measured event light curve, results from the combination of other lens parameters by \\begin{equation} t_{\\rm E}=\\left[ {4GM\\over c^2 v^2}{D_{ol}(D_{os}-D_{ol})\\over D_{os}} \\right]^{1/2}, \\end{equation} where $M$ is the lens mass, $v$ is the lens-source transverse speed, and $D_{ol}$ and $D_{os}$ represent respectively the distances to the lens and the source from the observer.  As a method to resolve the lens parameter degeneracy for general microlensing events, it was proposed to measure the shifts of the source star image centroid caused by lensing from follow-up astrometric observations for events detected photometrically from the ground by using high precision interferometers from space-based platform such as the {\\it Space Interferometry Mission} (SIM) and ground-based interferometers soon available on several 8 -- 10 m class telescopes such as the Keck and the Very Large Telescope (Miyamoto \\& Yoshii 1995; H\\o\\hskip-1pt g, Novikov, \\& Polarev 1995; Walker 1995; Dominik \\& Sahu 2000).   When the lens-source separation is normalized by the angular Einstein ring radius $\\theta_{\\rm E}$, the centroid shift vector with respect to the position of the un-lensed source for an event caused by a point-mass lens is represented by \\begin{equation} \\deltavec\\thetavec_{\\rm c} = {\\uvec \\over u^2+2} \\theta_{\\rm E}; \\end{equation} \\begin{equation} u_x = {t-t_0\\over t_{\\rm E}},\\ u_y = u_{\\rm min}, \\end{equation} where $\\uvec$ is the lens-source separation normalized by $\\theta_{\\rm E}$, $u_x$ and $u_y$ represent the components of $\\uvec$ parallel and normal to the lens proper motion vector $\\muvec$, $t_0$ is the time of the closest lens-source approach, and $u_{\\rm min}$ is the lens-source separation (also normalized by $\\theta_{\\rm E}$) at that moment (i.e.\\ the impact parameter). The angular Einstein ring radius is related to the physical parameters of the lens by \\begin{equation} \\theta_{\\rm E} = {vt_{\\rm E}\\over D_{ol}} = \\left( {4GM\\over c^2}{D_{os}-D_{ol}\\over D_{ol}D_{os}} \\right)^{1/2}. \\end{equation} When $u=\\sqrt{2}$, the amount of the centroid shift becomes maximum with \\begin{equation} \\delta\\theta_{\\rm c,max}= {\\theta_{\\rm E}\\over\\sqrt{8}}. \\end{equation} During the event, the trajectory of the centroid shifts traces an ellipse (astrometric ellipse).  Once the trajectory is constructed from astrometric observations of the event, one can determine $\\theta_{\\rm E}$ because the size of the astrometric ellipse (semi-major axis) is directly proportional to $\\theta_{\\rm E}$.  Then one can determine the lens proper motion by $\\mu=\\theta_{\\rm E}/t_{\\rm E}$ combined with the photometrically determined value of $t_{\\rm E}$.  With the determined value of $\\mu$, the uncertainty in the determined lens mass can be substantially reduced because $\\mu$ does not depend on $v$.  However, the lens parameter degeneracy is not completely broken even with the determination of $\\mu$, because it still results from the combination of the lens mass and the location.    The lens parameter degeneracy can be completely broken, and thus the lens mass can be uniquely determined if one determines the lens parallax by astrometrically measuring the deviations in the observed centroid shift trajectory, $\\Deltavec \\deltavec\\thetavec_{\\rm c}$, caused by the change of the observer's location during the event due to the orbital motion of the Earth around the Sun: parallax effect.\\footnote{The lens mass can also be determined by astrometrically monitoring nearby very high proper motion stars working as a lens instead of monitoring lensed source stars because the distance to the lens is known.  This method was suggested by Miralda-Escud\\'e (1996) and Paczy\\'nski (1998) and the specific candidate stars from the SIM monitoring were selected by Salim \\& Gould (2000).  We note, however, that our analysis is focused on followup observations for events detected by the photometric lensing surveys from the ground.} However, $\\Deltavec\\deltavec \\thetavec_{\\rm c}$, in general, is much smaller than $\\deltavec \\thetavec_{\\rm c}$, and thus measurements of the parallax-induced deviations require higher precision observations.  The possibility of astrometric parallax measurements was investigated by Paczy\\'nski (1998) and further analysis considering the actual performance of the specific instruments was done by Boden, Shao, \\& Van Buren (1998).  However, their analyses were based on several example events.   In this paper, we determine the distributions of the maximum centroid shifts, $f(\\delta\\theta_{\\rm c, max})$, and the average maximum deviations, $f(\\langle\\Delta\\delta\\theta_{\\rm c,max}\\rangle)$, for different types of Galactic microlensing events caused by various masses.  From the analysis of these distributions, we statistically investigate the detectabilities of $\\deltavec\\thetavec_{\\rm c}$ and $\\Deltavec\\deltavec\\thetavec_{\\rm c}$ expected from ground- and space-based observations.    \\begin{figure} \\epsfysize=10cm \\centerline{\\epsfbox{fig1.eps}} \\caption{ The parallax effect on the astrometric behavior of an example halo-LMC event.  The upper left panel shows the source star trajectories with respect to the lens ($L$) with (solid curve) and without (dotted curve) the parallax effect.  The dashed circle centered at $L$ represents the Einstein ring.  In the upper right panel, the trajectories of the centroid shifts $\\deltavec\\thetavec_{\\rm c}$ resulting from the individual source trajectories in the upper left panel are marked by the same line types. In the lower panel, we present the trajectory of the deviation vector $\\Deltavec\\deltavec\\thetavec_{\\rm c}$, which corresponds to the vector difference between the solid and dotted curve in the upper right panel. Both the trajectories of $\\deltavec\\thetavec_{\\rm c}$ and $\\Deltavec\\deltavec \\thetavec_{\\rm c}$ are presented in the ecliptic coordinates.  For the details about the lens system, see the text. } \\end{figure} ",
        "conclusions": "We compute the distributions of the maximum astrometric source star image centroid shifts, $\\delta\\theta_{\\rm c,max}$, and the average maximum deviations in the centroid shift trajectories, $\\langle\\Delta\\delta \\theta_{\\rm c,max}\\rangle$, expected for different types of Galactic events caused by various masses.  The findings from the analysis of these distributions are summarized as follows. \\begin{enumerate} \\item As long as source stars are bright enough for astrometric observations, one can detect $\\deltavec\\thetavec_{\\rm c}$ for most events caused by lenses with masses greater than $0.1\\ M_\\odot$ regardless of the event types from observations by using not only the SIM but also the ground interferometers.  \\item If observations are performed with the ground interferometers, detecting $\\Deltavec\\deltavec\\thetavec_{\\rm c}$ will be possible for nearly none of bulge self-lensing events and for only a small fractions of disk-bulge and halo-LMC events.  \\item However, if events are observed by using the SIM, detection of $\\Deltavec \\deltavec\\thetavec_{\\rm c}$ will be possible for most disk-bulge and halo-LMC events and even for some fraction of bulge self-lensing events. \\end{enumerate} Therefore, for the complete resolution of the lens degeneracy and accurate lens mass determination, SIM observations (or equivalent) will be essential."
    },
    "0011/astro-ph0011428_arXiv.txt": {
        "abstract": "We present high spatial resolution ($\\sim$0\\farcs6) near--infrared broad--band $JHK$ images and \\BG ~2.1661~$\\mu$m and \\H2 ~1--0 S(1) 2.122~$\\mu$m emission line images of the nuclear regions in the interacting starburst galaxies NGC 520, NGC 1614 and NGC 7714. The near--infrared emission line and radio morphologies are in general agreement, although there are differences in details. In NGC 1614, we detect a nuclear double structure in \\BG, in agreement with the radio double structure. We derive average extinctions of A$_K$ = 0.41 and A$_K$ = 0.18 toward the nuclear regions of NGC 1614 and NGC 7714, respectively. For NGC 520, the extinction is much higher, A$_K$ = 1.2 -- 1.6. The observed H$_2$/\\BG ~ratios indicate that the main excitation mechanism of the molecular gas is fluorescence by intense UV radiation from clusters of hot young stars, while shock excitation can be ruled out.  The starburst regions in all galaxies exhibit small \\BG ~equivalent widths. Assuming a constant star formation model, even with a lowered upper mass cutoff of M$_u$ = 30 M$_\\odot$, results in rather old ages (10 -- 40 Myr), in disagreement with the clumpy near--infrared morphologies. We prefer a model of an instantaneous burst of star formation with M$_u$ = 100 M$_\\odot$, occurring $\\sim$6--7 Myr ago, in agreement with previous determinations and with the detection of W--R features in NGC 1614 and NGC 7714. Finally, we note a possible systematic difference in the amount of hot molecular gas between starburst and Seyfert galaxies. ",
        "introduction": "Interaction and merging play an important role in the evolution of galaxies. Interacting galaxies usually have strong infrared (IR) emission (e.g. Rieke et al. 1980; Telesco, Wolstencroft \\& Done 1988), and show a wide range of nuclear activity, including AGN, nuclear starbursts (SB), ultraluminous IR galaxies and post--SB activity. These phenomena probably all arise from the ability of interactions to transport gas into the galactic nuclei (e.g. Barnes \\& Hernquist 1991) to fuel both AGN and SBs. In SB galaxies, the radiation output is dominated by star formation (SF), and SBs typically have SF rates of 5--50 M$_\\odot$ yr$^{-1}$ within a region of 0.1--1 kpc extent, much larger than that found in the Galactic center (e.g. Mezger, Duschl \\& Zylka 1996) or in normal galaxies (e.g. Keel 1983), and they thus provide an extreme environment in which to study the SF process.  Multiwavelength studies of interacting SB galaxies can provide insights into the connection between interactions, massive star clusters, and nuclear activity. The relatively unobscured near--IR (NIR) emission provides a better handle than optical emission to quantify the properties of SF in galaxies. In this paper we present high resolution NIR broad--band $JHK$ images and \\BG ~and \\H2 ~1--0 S(1) emission line images of the circumnuclear regions of three interacting/merging SB galaxies, NGC 520, NGC 1614 and NGC 7714. Because of their proximity, brightness and reasonably large angular size, they provide excellent targets for these studies. \\BG ~originates from H II regions surrounding hot young OB star clusters, while \\H2 ~arises from hot molecular gas and traces the material available for SF.  In the remaining part of Section 1, we give brief introduction to the galaxies. In Section 2 we discuss the observations and data reduction. In Section 3 we discuss the morphology of the circumnuclear regions, determine the extinctions to the SF complexes, constrain their SF properties, stellar populations and SF history by comparison among the NIR tracers and with multiwavelength emission, and discuss the gas masses of the galaxies. Conclusions are drawn in Section 4.  \\subsection{NGC 520}  NGC 520 (Arp 157, UGC 966, VV 231) at $v_\\mathrm{sys}$ = 2059 km s$^{-1}$ ($D = 27.3$ Mpc, 1$''$ = 130 pc, for H$_0=$ 75 Mpc$^{-1}$ km s$^{-1}$ and q$_0$ = 0.5) is an IR--luminous peculiar pair of galaxies with a highly disturbed flattened morphology (e.g. Stanford \\& Balcells 1990; Hibbard \\& van Gorkom 1996). The inclination of NGC 520 is $\\sim$66\\degr ~(Rownd \\& Young 1999). NGC 520 is at an intermediate merger stage (Hibbard \\& van Gorkom 1996), with the two nuclei embedded within a common stellar envelope. The primary nucleus (hereafter PN) is optically hidden behind a dust lane at PA = 95\\degr, while the second, optically bright nucleus (hereafter NWN) is situated $\\sim$40$''$ (5.3 kpc) to NW. A bright optical tail stretches $\\sim$2\\farcm8 (22 kpc) to SE, bending sharply to E and connecting to a broad stellar plume.  The PN has $\\sim$4.3 $\\times$ 10$^9$ M$_\\odot$ of molecular gas in a $\\sim$7$''$ $\\times$ 3$''$ (920 $\\times$ 400 pc) scale rotating molecular gas disk at PA = 95\\degr ~(Yun \\& Hibbard 2000). While the CO emission is in good agreement with the 1.4 GHz radio emission (5$''$ $\\times$ 2$''$ = 660 $\\times$ 260 pc, PA = 93\\degr; Condon et al. 1990), both avoid the \\HA ~emission, which is dominated by plumes of ionized gas emerging up to a projected distance of 20$''$ (2.6 kpc) from the nucleus along PA = 25\\degr ~(Hibbard \\& van Gorkom 1996). These plumes probably represent a SB--driven bipolar outflow of ionized gas out of the nuclear SB region (Norman et al. 1996). While \\HA ~emission regions exist within the PN region, the {\\em nuclear} SB is completely obscured optically by dust within the nuclear gas disk (but becomes visible in \\BG ~; Section 3.1.1).  NGC 520 is probably the result of an encounter $\\sim$3 $\\times$ 10$^8$ yr ago between a gas--rich and a gas--poor galaxy (Stanford \\& Barcells 1991). Both the S tail and the plume arose from the gas--poor progenitor and are currently disturbing the underlying gas distribution from the gas--rich disk (Hibbard \\& van Gorkom 1996). The optical spectrum of the NWN is dominated by A stars, indicating that it is now in a post-SB phase (Stanford \\& Balcells 1990; Bernl\\\"ohr 1993b), in agreement with the non--detection of \\BG ~emission (Section 3.1.1). Enhanced SF is currently detected only in the PN (e.g. Stanford 1991).  \\subsection{NGC 1614}  NGC 1614 (Mrk 617, Arp 186, II Zw 15) is an IR--luminous SB(s)c pec type SB galaxy at $v_\\mathrm{sys}$ = 4723 km s$^{-1}$ ($D = 62.4$ Mpc, 1$''$ = 300 pc). It has been extensively studied at optical (e.g. De Robertis \\& Shaw 1988), NIR (e.g. Aitken, Roche \\& Phillips 1981; Forbes et al. 1992; Puxley \\& Brand 1999) and radio (e.g. Condon et al. 1982) wavelengths. The inclination of NGC 1614 is $\\sim$30\\degr ~(Rownd \\& Young 1999).  The extended emission around NGC 1614 is highly asymmetric, suggesting recent tidal interaction with at least one other galaxy and that the galaxies are currently merging. The chaotic optical structure includes two roughly symmetrical inner spiral arms, a linear tail to the SW extending almost 1$'$ (18 kpc) from the nucleus and a large curving arc to the E/SE that reaches $\\sim$33$''$ (9.8 kpc) from the nucleus (e.g. Neff et al. 1990). Only a single peak is seen in the NIR continuum, suggesting that the two galaxies have already merged. However, two peaks (possibly double nuclei) are seen in the \\BG ~emission (Section 3.1.2) and in radio (Condon et al. 1982).  Armus, Heckman \\& Miley (1989) found a weak broad feature at 4660 \\AA ~due to W--R stars in the nuclear spectrum, which was later confirmed by Vacca \\& Conti (1992). Neff et al. (1990) made a multiwavelength study of NGC 1614 and concluded that its extreme IR luminosity results from vigorous nuclear SF induced by the interaction and merger of at least two galaxies. The molecular hydrogen mass of NGC 1614 is M(H$_2$) = 1.1 $\\times$ 10$^{10}$ M$_\\odot$ (Sanders, Scoville \\& Soifer 1991). There is an unresolved central CO source in NGC 1614 which emits $\\sim$30 \\% of the total CO flux (Scoville et al. 1989), implying a very high molecular gas concentration to the nucleus.  \\subsection{NGC 7714}  NGC 7714 (Arp 284, VV 51, Mrk 538) is a nearby ($v_\\mathrm{sys}$ = 2808 km s$^{-1}$, $D = 37.2$ Mpc, 1$''$ = 180 pc) peculiar SBb type SB galaxy. It is very IR--luminous and has strong SB activity both in the nucleus and in circumnuclear regions (e.g. Gonzalez-Delgado et al. 1995, hereafter GD95; Garcia-Vargas et al. 1997). NGC 7714 and NGC 7715 (at $\\sim$2$'$ (21 kpc) to the E; $v_\\mathrm{sys}$ = 2770 km s$^{-1}$, $D = 36.7$ Mpc) form the Arp 284 system of interacting spirals, as is evident from the HI and optical bridge (Smith, Struck \\& Pogge 1997; Papaderos \\& Fricke 1998) between them. The inclination of NGC 7714 is $\\sim$42\\degr ~(Chapelon, Contini \\& Davoust 1999). A striking feature of NGC 7714 is an asymmetric circumnuclear ring of diameter $\\sim$20$''$ $\\times$ 40$''$ (3.6 $\\times$ 7.2 kpc), with most of the emission located to the E of the nucleus. This ring is largely devoid of \\HA ~(GD95) and HI emission (Smith et al. 1997), and has NIR colours similar to the old disk population (Bushouse \\& Werner 1990). This structure is thus probably related to the dynamical perturbation of the disk rather than to previous SF. A weak He II 4686 \\AA ~W--R feature has been detected in NGC 7714 (GD95). The molecular gas mass of NGC 7714 is M(H$_2$) = 2.1 $\\times$ 10$^9$ M$_\\odot$ (Sanders et al. 1991).  While NGC 7714 shows strong SB activity, NGC 7715 has no optical emission lines and is in a post--SB phase (Bernl\\\"ohr 1993a). There are three main circumnuclear HII regions in NGC 7714, labelled A, B and C by GD95, and located in the bulge/disk of the galaxy at 5$''$ (900 pc) E, 12$''$ (2.1 kpc) NW and 14$''$ (2.5 kpc) SW from the nucleus, respectively. Region A is detected by us as \\BG ~region 4. In the deconvolved \\HA ~image of GD95, the nuclear SB region breaks furthermore into the nucleus and an extranuclear component, which we identify as \\BG ~region 2. ",
        "conclusions": "We present high spatial resolution ($\\sim$0\\farcs6) near--infrared broad--band $JHK$ images and \\BG ~2.1661~$\\mu$m and \\H2 ~1--0 S(1) 2.122~$\\mu$m emission line images of the nuclear regions in the interacting starburst galaxies NGC 520, NGC 1614 and NGC 7714. The near--infrared emission line and radio morphologies are in general agreement, although there are differences in details. In NGC 1614, we detect a nuclear double structure in \\BG, in agreement with the radio double structure. We derive average extinctions of A$_K$ = 0.41 and A$_K$ = 0.18 toward the nuclear regions of NGC 1614 and NGC 7714, respectively. For NGC 520, the extinction is much higher, A$_K$ = 1.2 -- 1.6. The observed H$_2$/\\BG ~ratios indicate that the main excitation mechanism of the molecular gas is fluorescence by intense UV radiation from clusters of hot young stars, while shock excitation can be ruled out.  The starburst regions in all galaxies exhibit small \\BG ~equivalent widths. Assuming a constant star formation model, even with a lowered upper mass cutoff of M$_u$ = 30 M$_\\odot$, results in rather old ages (10 -- 40 Myr), in disagreement with the clumpy near--infrared morphologies. We prefer a model of an instantaneous burst of star formation with M$_u$ = 100 M$_\\odot$ occurring $\\sim$6--7 Myr ago, in agreement with previous determinations and with the detection of W--R features in NGC 1614 and NGC 7714. Finally, we note a possible systematic difference in the amount of hot molecular gas between starburst and Seyfert galaxies."
    },
    "0011/astro-ph0011334_arXiv.txt": {
        "abstract": "s{Observational evidence for the presence of a preferred scale around 100~Mpc in the distribution of high-density regions in the Universe is summarised.  Toy models with various degrees of regularity are analysed to understand better the observational data. Predictions of models which produce a feature on the initial power spectrum of matter are analysed and compared with new observational data on the CMB angular spectrum, cluster mass function, and galaxy and cluster power spectrum.  } ",
        "introduction": "Until the mid-1970s the astronomical community accepted the paradigm that galaxies are distributed in space more or less randomly and that about $5 - 10$~\\% of all galaxies are located in clusters which are also randomly distributed. A paradigm change occurred during a Symposium on Large-Scale Structure of the Universe in 1977 in Tallinn. Several teams reported results of studies of the 3-dimensional distribution of galaxies in superclusters, and demonstrated that superclusters consist of chains (filaments) of galaxies, groups and clusters, and that the space between galaxy systems is devoid of any visible galaxies (J\\~oeveer \\& Einasto 1978, Tarenghi \\etal 1978, Tifft \\& Gregory 1978, Tully \\& Fisher 1978).   The discovery of the filamentary character of superclusters and the presence of large empty voids made obsolete the early theories of structure formation -- the hierarchical clustering model by Peebles (1980), and the whirl theory by Ozernoy (1978). The first impression was that the winning theory was the pancake scenario by Zeldovich (1978).  However, a more detailed analysis has shown that the original pancake scenario also has weak points -- the structure forms too late and has no fine structure of faint galaxy filaments in large voids (Zeldovich, Einasto \\& Shandarin 1982).  Then the Cold Dark Matter model was suggested and generally accepted (Blumenthal et al 1984). Recently it was replaced by a CDM model dominated by a cosmological term (vacuum energy) -- the LCDM model.  However, this model has also at least one weak point: in LCDM model superclusters are randomly distributed, whereas observational evidence suggests that rich superclusters and voids form a quasi-regular network of scale $\\sim 100 - 130$~\\Mpc\\ (Einasto \\etal 1997a, 1997d). Here we use the dimensionless Hubble constant $h$ defined as $H_0 = 100~h$ km~s$^{-1}$~Mpc$^{-1}$.  These data raise the question: Is the regularity of the distribution of superclusters a challenge for CDM models?  My lecture tries to answer this question. First I shall summarise the observational evidence for the presence of the 130~\\Mpc\\ scale, thereafter I shall discuss attempts to explain the presence of this scale theoretically. ",
        "conclusions": "\\begin{itemize} \\vspace*{-2mm} \\item{} There exists good evidence on the presence of a scale of $\\sim 130$~\\Mpc: it measures the mean distance between rich superclusters across voids. \\vspace*{-2mm} \\item{} There are variants of the inflation scenario suggesting the presence of a characteristic scale (bump) in the power spectrum. \\vspace*{-2mm} \\item{}  Models with a bump in the power spectrum are in agreement with other cosmological data using sets of parameters within acceptable ranges.  \\vspace*{-2mm} \\item{} Cellular large-scale structure may be the end of the fractal structure of the Universe, as suggested by Ruffini, Song \\& Taraglio (1988). \\end{itemize}  We should note that models neither predict the position of the bump, nor the presence and extent of the regularity of the supercluster-void network. We should remember the maxim by Eddington: {\\em no experimental result can be taken seriously if not explained theoretically}.  Thus larger surveys are needed to clarify properties of the supercluster-void network.  Also a further theoretical analysis are needed to find the physical origin of the scale and the extent of the regularity."
    },
    "0011/astro-ph0011044_arXiv.txt": {
        "abstract": "We present accurate analytic fitting formulae for the frequency integrated Gaunt factors for the relativistic as well as nonrelativistic thermal bremsstrahlung which is relevant to the intracluster plasma.  The fitting is carried out for $1 \\leq Z_j \\leq 28$, $6.0 \\leq \\log T \\leq 8.5$, where $Z_j$ is the charge number of the ion, and $T$ is the electron temperature in kelvins.  The present analytic fitting formulae will be useful for the analysis of the X-ray emission which comes from the intracluster plasma as well as the other X-ray sources. ",
        "introduction": "Nozawa, Itoh, \\& Kohyama (1998) have recently carried out accurate  calculations on the relativistic thermal bremsstrahlung Gaunt factor for the intracluster plasma.  Their calculation is based on the method of Itoh and his collaborators (Itoh, Nakagawa, \\& Kohyama 1985; Nakagawa, Kohyama, \\& Itoh 1987; Itoh, Kojo, \\& Nakagawa 1990; Itoh et al. 1991,1997). In calculating the relativistic thermal bremsstrahlung Gaunt factor for the high-temperature, low-density plasma, Nozawa, Itoh, \\& Kohyama (1998) have made use of the Bethe-Heitler cross section (Bethe \\& Heitler 1934) corrected by the Elwert factor (Elwert 1939). They have also calculated the Gaunt factor using the Coulomb-distorted wave functions for nonrelativistic electrons following the method of Karzas \\& Latter (1961). They have presented the results of the numerical calculations in the form of extensive numerical tables. The grid of their numerical table is made fine enough so that smooth interpolations can be carried out for the applications of the results to the analyses of various  astrophysical observational data.  Other references on the calculation of thermal bremsstrahlung Gaunt factor include Culhane (1969), Culhane \\& Acton (1970), Raymond \\& Smith (1977), Gronenschild \\& Mewe (1978), Mewe, Lemen, \\& van den Oord (1986), and Carson (1988).  Itoh et al. (2000) presented accurate analytic fitting formulae which summarize the results of the extensive numerical tables presented by Nozawa, Itoh, \\& Kohyama (1998).  Analytic fitting formulae can be readily implemented in the computer programs for the analyses of the observational data. Therefore, we recommend these analytic fitting formulae be used widely by the observers for the analyses of their observational data.  In this paper we will present accurate analytic fitting formulae for the frequency integrated emissivity due to the relativistic as well as nonrelativistic thermal bremsstrahlung which is relevant to the intracluster plasma.  The fitting is carried out for $1 \\leq Z_j \\leq 28$, $6.0 \\leq \\log T \\leq 8.5$, where $Z_j$ is the charge number of the ion, and $T$ is the electron temperature in kelvins.  The present analytic fitting formula will be useful for the analysis of the X-ray emission which comes from the intracluster plasma as well as the other X-ray sources.  The present paper is organized as follows. We will give the analytic fitting formulae for the frequency integrated Gaunt factor for the relativistic thermal bremsstrahlung in $\\S$2.  We will give the analytic formula for the frequency integrated Gaunt factor for the nonrelativistic thermal bremsstrahlung in $\\S$3.  Concluding remarks will be given in $\\S$4. ",
        "conclusions": "We have presented accurate analytic fitting formulae for the frequency integrated Gaunt factors for relativistic as well as nonrelativistic thermal bremsstrahlung.  The analytic fitting formulae have been constructed from the numerical results of the calculation reported in Nozawa, Itoh, \\& Kohyama (1998).  The accuracy of the fitting is generally better than 0.1\\%. The present fitting formulae can be used widely for the analysis of thermal bremsstrahlung radiation.  We thank Professor Y. Oyanagi for allowing us to use the least square fitting program SALS.  This work is financially supported in part by the Grant-in-Aid of Japanese Ministry of Education, Science, Sports, and Culture under the contract \\#10640289.  \\newpage"
    },
    "0011/astro-ph0011272_arXiv.txt": {
        "abstract": "{ CCD photometry of the classical nova RW~UMi during 10 nights in 1995 through an $R$ filter and 14 nights in 1997 through a `clear' filter, with the 1-m Wise telescope yields a periodic modulation in the light curve with the period 0.05912$\\pm$0.00015 d and a semi amplitude of 0.025 mag. We discuss several explanations for the periodicity, and suggest that it is likely to be the orbital period of the binary system or a superhump period. RW~UMi is thus a candidate for the classical nova with the shortest known orbital period. } ",
        "introduction": "RW~UMi is a high galactic latitude nova that erupted in 1956. Its outburst was reported, however, only about six years later (Kukarkin 1962). The nova reached a maximum of at least $m_{V}$$\\approx$$6.0$ on 1956 September 24, and decayed to $m_{V}$$\\approx$11.5 a year later. Kukarkin also concluded that the progenitor of the nova was fainter than $m_{V}$=21.2 in POSS. Duerbeck (1987), however, barely detected the pre-nova at $m_{V}$=21.0 in the same plate. Photometric observations of the object during the last two decades are consistent with a visual magnitude stable around $m_{V}$$\\approx$18.8 (Cohen 1985; Kaluzny \\& Chlebowski 1989; Szkody et al. 1989; Howell et al. 1991; Szkody \\& Howell 1992; Szkody 1994; Ringwald et al. 1996; Downes \\& Duerbeck 2000). We note that the estimate of the nova in 1988 as $m_{V}$=21.0-21.5 (Howell 1988) was corrected by Szkody et al. (1989).  Spectra of the post-nova are presented by Kaluzny \\& Chlebowski (1989), Szkody \\& Howell (1992) and Ringwald et al. (1996). These show a strong blue continuum and several weak emission lines -- the Balmer and He II 4686 lines.    Cohen (1985) marginally detected a nebula around the remnant of RW~UMi using an image taken through a narrow H$\\alpha$ filter. Slavin et al. (1995) repeated this observation, and concluded that the distance to the binary system is 5$\\pm$2 Kpc. Esenoglu et al. (2000) resolved the shell of RW~UMi in 1995 and derived a distance range for the nova of 4200-5840 pc. Downes \\& Duerbeck (2000) failed, however, to detect the nova nebula in an observation carried out in 1998.  Kaluzny \\& Chlebowski (1989) found that the nova was not variable at the 0.05 mag level in the $B$ and $V$ filters in five images taken in 1985 and one in 1987. Szkody et al. (1989) reported the discovery of a variation of about 0.2 mag in a 3-h observation in 1988 using a wide band red filter. They suggested a possible period of 117$\\pm$5 minutes. Shafter \\& Campbell (1990) independently announced a similar detection. Howell et al. (1991) observed the nova with a $B$ filter for 2.1 h in 1989, and proposeed a period of 113$\\pm$10 minutes.  There are only three recorded classical novae with established orbital periods below the cataclysmic variables (CVs) period gap -- CP~Pup 1942, GQ~Mus 1983 and V1974~Cyg 1992 (Diaz \\& Bruch 1997; Ritter \\& Kolb 1998\\footnote{Note that the dwarf novae WX~Cet and TV~Crv were misclassified as novae.}). This is compared with about one hundred and twenty other short orbital period systems. In comparison, above the gap there are more than thirty novae and about one hundred and thirty other systems. Thus the ratio of novae to all CVs below the gap is nearly ten times smaller than above the gap. An increase of the number of known novae below the period gap is, therefore, necessary for the understanding of nova cycles of short orbital period CVs. The possibility that the orbital period in RW~UMi is below the gap lead us to add this nova to our list of targets in an ongoing programme, maintained at the Wise Observatory, to search for periods in post-novae. Retter (1999) announced the detection of a 1.4-h periodicity in the light curve of RW~UMi. In this work we elaborate that report, discuss possible physical mechanisms for the variation and further try to classify the system within the known subgroups of CVs. ",
        "conclusions": "The light curve of RW~UMi in 1997 was modulated by the 1.4-h period. There are three obvious interpretations for this periodicity. It might represent the spin period of the white dwarf, a superhump variation or the orbital period of the binary system. In the following we examine the different models and try to classify the system in the light of the observations accumulated so far on the nova.  \\subsection{A spin period-?}  First, we discuss the question whether the system is an AM~Her (polar) system (for a review see Warner 1995). The primary white dwarfs in these systems possess strong magnetic fields. Thus, the accreted matter from the companion star flows directly to the two poles of the rotating white dwarf whose spin period is synchronized with the orbital period of the binary system.   The spectra of RW~UMi show a strong continuum and only weak emission lines. There are no strong high excitation lines, and the He II 4686 line is rather weak (Sect.~1). We thus believe that the polar scenario can be eliminated, since typical spectra of AM~Her systems have a weak continuum and strong emission lines (Warner 1995). The typical optical spectra, radiated from accretion discs in post-novae, are characterized by a strong continuum, because the disc is usually optically thick in classical nova systems for a few decades after their eruptions (Retter et al. 1999).  If the magnetic field is moderate the spin period is usually much shorter than the orbital period and an accretion disc might be formed around the white dwarf. These objects are termed `intermediate polars' (for reviews see Patterson 1994; Hellier 1996). Spin periods in intermediate polars range between 33 s in AE Aqr (Hellier 1996) to 1.44 h in Nova V1425~Aql 1995 (Retter et al. 1998). RW~UMi might thus be an intermediate polar with the 1.4-h period being the rotation period of the white dwarf. The orbital period should then be somewhat longer than the spin period. The power spectrum of RW~UMi (Fig.~3) does not show, however, other periodicities in addition to the 1.4-h period (Sect.~3). Therefore, it seems unlikely that the periodicity we have discovered is a spin period.   \\begin{table*} \\label{table:v1500refs} \\caption[]{Properties of the classical novae below the period gap} \\begin{center} \\begin{tabular}{lcllccc}  Object name&year of outburst&orbital period& nature         &pre-nova magnitude&post-nova magnitude&difference\\\\  \\\\  RW~UMi     & 1956   & 1.419-?$^{1}$  & non-magnetic-?$^{1}$ & $21^{2}$         & 18.8$^{1}$        & 2.2      \\\\ GQ~Mus     & 1983   & 1.425$^{3,4}$  & magnetic-?$^{3,4}$   & $22^{2}$         & 18.0-?$^{4,5}$    & 4-?      \\\\ CP~Pup     & 1942   & 1.474$^{6}$    & non-magnetic$^{6}$   & $>$$17^{7}$        & 15.3$^{6}$        & $>$$1.7$   \\\\ V1974~Cyg  & 1992   & 1.950$^{8,9}$  & non-magnetic$^{8,9}$ & 21$^{8,10}$      & 16.5--?$^{11}$    & 4.5--?   \\\\  \\end{tabular} \\end{center}   $^1$This work; $^2$Duerbeck 1987; $^3$Diaz \\& Steiner 1989; $^4$Diaz \\& Steiner 1994; $^5$Diaz et al. 1995; $^6$Patterson \\& Warner 1998; $^7$Warner 1995; $^8$Retter \\& Leibowitz 1998; $^9$Skillman et al. 1997; $^{10}$Pavelin et al. 1993; $^{11}$Goransky 2000.  \\end{table*}   \\subsection{A superhump period-?}  Superhumps are quasi-periodic oscillations that have been observed in the light curves of non-magnetic CVs with relatively short orbital periods (typically $P_{orbital}$$<$4 h). They either appear in the bright outbursts (superoutbursts) of the SU~UMa subclass of dwarf novae or permanently in nova-likes. Superhump periods differ from binary periods by a few percent. They are understood as the beat periods of the orbital period and the apsidal (or the nodal) precession period of the accretion disc. For reviews on superhumps see Warner (1995) and Osaki (1996) and on permanent superhumps -- Patterson (1999) and Retter \\& Naylor (2000). If the 1.42-h period in RW~UMi is a positive superhump, a superhump period excess of 1-4\\% is expected according to the observed values in superhump systems (Ritter \\& Kolb 1998), leading to a binary period in the range 1.36-1.40 h. If the periodicity is a negative superhump, the expected range in the binary period is 1.43-1.45 h since period deficits in negative superhumps are about half period excesses in positive superhumps (Patterson 1999). In both cases the orbital period is only a few percent different from the observed periodicity.  \\subsection{The binary period-?}  Orbital periods of nova systems range between 1.425 h (GQ~Mus) to 48 h (GK~Per) with a peak around 3-4 h (Diaz \\& Bruch 1997). The observed orbital period distribution of CVs presents a deficiency of systems between about 2 h and 3 h, which is termed 'the period gap' (Warner 1995). So far only three novae with established orbital periods below the gap have been discovered. This is compared with a few dozen novae with orbital periods above the gap (Sect.~1). If the 1.419-h periodicity observed in the light curve of RW~UMi is its orbital period, it would be the shortest recorded binary period in a classical nova, although confirmation by radial velocity measurements is required. Table~1 presents the pre-nova and post-outburst magnitudes of the novae below the gap as well as several other properties. The data show that the brightness of the post-novae below the gap is still above the pre-nova state. RW~UMi is much brighter than the progenitor value as well.    \\subsection{RW~UMi as a high galactic latitude nova}  RW~UMi is a high galactic latitude nova and therefore probably a binary system in the halo of the galaxy. Halo CVs are believed to be relatively old systems compared with the population of CVs in the galactic disc (see e.g. Howell \\& Szkody 1990). Thus, according to evolutionary theories halo systems should have shorter orbital periods (assuming they haven't bounced yet at the period minimum). Such a difference between the two populations has indeed been found (Howell \\& Szkody 1990). The possibility that the orbital period of RW~UMi is the shortest among all known novae is consistent with this idea.  \\subsection{The mass accretion rate and thermal stability of RW~UMi}  In this section we estimate the mass transfer rate in RW~UMi to assess whether it has dwarf nova or nova-like characteristics. Permanent superhumpers (and nova-likes) are believed to occur in CVs with thermally stable discs, while dwarf novae (including SU~UMa systems) are thermally unstable. Based on the thermal-tidal disc instability model (Osaki 1996), assuming that the disc is the dominant light source in the visual band, and neglecting a possible contribution from the hot white dwarf in post-novae, Retter \\& Leibowitz (1998) and Retter \\& Naylor (2000) developed a way to check the thermal stability state of disc-accreting CVs. In the following, we apply this method to RW~UMi to obtain its mass transfer rate.  The visual magnitude of the post-nova is $m_{V}$$\\approx$18.8 and the distance -- 4200-5840 pc (Sect.~1). The interstellar reddening is $A_{V}$=0.1 (Anhert 1963; Cohen 1985). For the mass of the white dwarf, we take $M_{WD}$=0.7$M_{\\odot}$ -- typical to CVs below the gap (Patterson 1998). The absence of eclipses (Sect.~3) limits the inclination angle to i=0-65$^{\\circ}$. Inserting all these values into eq.(3) of Retter \\& Naylor (2000), we obtain $\\dot{M}$=1.2-8.6$\\times 10^{17}$g s$^{-1}$. The critical mass transfer rate, calculated by eq.(2) of Retter \\& Leibowitz (1998), assuming that the orbital period is 1.4 h (Sect.~4.3) is 0.5$\\times 10^{17}$g s$^{-1}$. The system is thus very likely to be thermally stable. This fact together with the long-term light curve of the nova (Fig.~1) that does not show dwarf nova outbursts (although obviously outbursts could have been missed) seem to reject the possibility that the post-nova is a dwarf nova system, and therefore it should be a nova-like system. A similar test applied to the progenitor of the nova at $m_{V}$=21.0 (Sect.~1) yields $\\dot{M}$=0.2-1.1$\\times 10^{17}$g s$^{-1}$. This range straddles the limiting value, so the thermal stability state of the pre-nova cannot be determined.  \\subsection{Final classification of the system}  A definite classification of the system cannot be made. An AM~Her model doesn't fit the nova properties (Sect.~4.1). An intermediate polar model might be applicable to RW~UMi, but only if an additional short-term periodicity (the spin period) or a relatively long period (which would be the orbital period, making the 1.4-h period the spin period) is found in the future. We note that it is very likely that an accretion disc is formed around the white dwarf in RW~UMi as the AM~Her scenario has been rejected and since it is believed that most intermediate polars have accretion discs (Patterson 1994).  A dwarf nova (SU~UMa) model seems to be ruled out by the long-term behaviour of the light curve in 1995-1999 (Fig.~1) that does not show any dwarf nova outburst, and by our findings that the system is almost certainly thermally stable (previous section). RW~UMi is therefore most likely a system with a short orbital period in a high state. As such, it should have permanent superhumps according to Osaki (1996), however, the observations accumulated so far cannot confirm this phenomenon, probably because the nova is relatively faint."
    },
    "0011/astro-ph0011507.txt": {
        "abstract": "Recent shell-model calculations of weak-interaction rates for nuclei in the mass range $A = 45-65$ have resulted in substantial revisions to the hitherto standard set of Fuller, Fowler, \\& Newman (FFN).  In particular, key electron-capture rates, such as that for \\I{60}{Co} are much smaller.  We consider here the effects of these revised rates on the presupernova (post-oxygen burning) evolution of massive stars in the mass range 11 to 40\\,\\Msun.  Moreover, we include, for the first time in models by our group, the effects of modern rates for beta-decay in addition to electron capture and positron emission. Stars of 15, 25, and 40\\,\\Msun stars are examined in detail using both the full FFN rate set and the revised rates.  An additional finely spaced (in mass) grid of 34 models is also calculated in order to give the systematics of iron core and silicon core masses. Values for the central electron mole number at the time of iron core collapse in the new models are typically larger, by $\\Delta \\Ye = 0.005$ to 0.015, than those of \\cite{WW95}, with a tendency for the more massive models to display larger differences. About half of this change is a consequence of including beta-decay; the other half, result of the smaller rates for electron capture. Unlike what might be expected solely on basis of the larger \\Ye's, the new iron core masses are systematically smaller owing to a decrease in the entropy in the outer iron core.  The changes in iron core mass range from zero to 0.1\\,\\Msun (larger changes for high mass stars). It would be erroneous though to estimate the facility of exploding these new models based solely upon their iron core mass since the entire core structure is altered and the density change is not so different as the adjustments in composition might suggest. We also observe, as predicted by \\cite{auf94a}, a tendency towards beta-equilibrium just prior to the collapse of the core, and the subsequent loss of that equilibrium as core collapse proceeds.  This tendency is more pronounced in the 15\\,\\Msun model than in the heavier stars. We discuss the key weak reaction rates, both beta-decay and electron-capture, responsible for the evolution of \\Ye and make suggestions for future measurements. ",
        "introduction": "\\lSect{intro}  The effect of weak interactions on the final evolution of massive stars is known to be quite pronounced \\cite[e.g.,][]{WFW84}.  Weak interactions alter the composition and change the nucleosynthesis. They modify the structure of the iron core and its overlying mantle and thus help to determine whether the star explodes and the characteristics of the bound remnant. Electron capture decreases the number of electrons available for pressure support and, all else being equal, decreases the effective Chandrasekhar Mass that can be supported without collapsing. Beta decay acts in the opposite direction. Weak interactions also create neutrinos that, for all densities until the collapse becomes truly hydrodynamic (i.e., $\\rho \\ltaprx 10^{11}$ g cm$^{-3}$), escape the star carrying away energy and thereby reducing entropy.  In fact, for stars that are not too massive, neutrinos from weak interactions, as opposed to thermal plasma processes, are the dominant energy sink in the presupernova star \\cite[]{WZW78}.  Reducing the entropy also makes the iron core smaller \\cite[]{Timmes96}.  Over the years, many people have calculated weak interaction rates for astrophysical application \\cite[]{Hans66,Hans68,Mazu73,MTC74,Taka73,auf94b}.  For approximately 15 years though, the standard in the field has been the tabulations of \\cite{FFN80,FFN82a,FFN82b,FFN85} (henceforth FFN).  These authors calculated rates for electron capture, beta-decay, and positron emission, plus the associated neutrino losses for all astrophysically relevant nuclei ranging in atomic mass from 23 to 60. Their calculations were based upon an examination of all available information (in the mid 1980's) for individual ground states and low-lying excited states in the nuclei of interest.  Recognizing that this only saturated a small part of the Gamow-Teller resonance, they extrapolated the known experimental rates using a simple (single-state) representation of this resonance.  Now, as we briefly summarize in \\Sect{newrates}, new shell models of the distribution of Gamow-Teller strength have resulted in an improved - and often reduced - estimate of its strength.  The alteration in some important rates, such as electron capture on $^{60}$Co, is as large as two orders of magnitude in important astrophysical circumstances.  In order to determine how these new rates affect the post-oxygen burning evolution of massive stars and presupernova models, we have examined the evolution of a fine grid of 34 stellar models spanning a mass range 11 to 40\\,\\Msun.  The post-oxygen burning evolution of each star was calculated twice - once using the same old set of weak rates as employed in \\cite{WW95} (henceforth WW) by Tom Weaver in 1996 (priv. com.; unpublished) and, again, with the new rates.  It is important to note that the models of WW included electron capture and positron decay rates from FFN, but used beta-decay rates from \\citeauthor{Hans68} and \\citeauthor{Mazu73}. These were much smaller than their FFN counterparts and essentially resulted in beta decay being neglected.  This was a mistake, that, so far as we are aware, was not replicated in the more recent works of other groups \\cite[]{Thie96,Chie98}.  An examination of conditions in the pre-supernova model led to the (correct) conclusion that beta-decay was negligible there. What was overlooked however, as was subsequently pointed out \\cite[]{auf94a}, was the importance of beta-decay in the iron core a few hours \\emph{before} iron core collapse.  The new weak rate implementation described here \\emph{does} include beta-decay - along with revised rates for electron capture and positron emission. In order to make a fair comparison to what would have been obtained by WW had they used the full FFN implementation, we have recalculated (\\Sect{stars}) three stars (15, 25, and 40\\,\\Msun) using FFN beta decay rates as well as FFN electron-capture and positron-decay rates keeping the stellar physics (thermal neutrino loss rates, opacities, mass loss, etc.) the same as in WW.  Subsequent survey of massive stellar evolution \\cite[]{Raus00,Hege01} will implement the full revised weak rate set along with many other improvements in a fine grid of masses and metallicities.  We find (\\Sect{newresults}) significant differences in the presupernova models calculated using the new rates In particular, the large decrease in key rates for electron-capture and the inclusion of beta-decay result in an electron mole number, \\Ye, in the center of the presupernova star that is 1\\,\\% to 4\\,\\% larger. Changes around 3\\,\\% are most common and there is a tendency for the biggest changes to occur for larger stellar masses.  The altered weak rates also affect the structure of the presupernova star, leading to iron core masses that are, almost without exception, \\emph{smaller}.  This seemingly paradoxical result (smaller Fe core and larger \\Ye) comes about because of a decreased entropy in the outer iron core, a consequence of a final stage of silicon shell burning.  In \\Sect{conclude} we summarize our findings and enumerate key rates that warrant further laboratory study (see also \\Tabs{15} -- \\Tabff{flow40}). ",
        "conclusions": "\\lSect{conclude}  We have explored in some detail the weak interactions that characterize the last few days in the life of a massive star. In particular we have studied the effects of a new, improved set of weak interaction rates and neutrino losses by \\cite{LM00} and the role played by beta-decay as well as electron capture.  In agreement with \\cite{auf94a}, we find that inclusion of beta decay makes a significant difference in the presupernova structure. Beta equilibrium comes to exist, temporarily, in the center of the star after silicon is depleted.  This causes the final electron mole number, \\Ye, to be larger than in models where beta decay has been neglected \\cite[e.g.,][]{WW95}.  The existence of beta equilibrium, particularly in the lower mass stars ($M \\ltaprx 20\\,\\Msun$), also decreases the central entropy. Both effects alter the structure of the presupernova star.  In addition to the consequences of including beta decay, we find that which rate set one uses also makes a difference. In particular, for the more massive stars studied here, use of the new weak rates of LMP instead of the older set from FFN causes an additional increase in \\Ye comparable to that from the inclusion of beta decay.  This is because the new rates for electron capture are systematically smaller owing to a more accurate treatment of the Gamow-Teller resonance. All in all, \\Ye in the iron cores of the new models is systematically larger by approximately 3\\,\\% (\\Fig{yefeofm}). Since \\Ye enters quadratically into determining the Chandrasekhar Mass and since the mass of the iron core is critically important in models for the supernova explosion mechanism, this is a significant difference.  However, contrary to simple expectations based only upon the change in \\Ye, the new iron core masses are systematically about 0.05 to 0.1\\,\\Msun \\emph{smaller} than in WW (\\Fig{yefeofm}).  This has several causes, but chiefly reflects the lower entropy in the convective silicon shell just prior to iron core collapse.  The lower entropy inhibits the final growth of the iron core in a surrounding medium of fixed entropy.  In total, the new presupernova models may explode more easier due both to their reduced iron core mass and larger \\Ye, however the changes are not large and their are other effects (lower entropy just outside the iron core) that could act in the opposite direction.  The new 15, 20, and 25\\,\\Msun models discussed here are available to those wanting to study the explosion mechanism.  As \\Tabs{15} -- \\Tabff{flow40} show, some of the most interesting weak processes in presupernova stars involve unstable, radioactive nuclei for which no experimental data exists about the GT strength distribution.  This situation should change once the proposed Rare-Isotope Accelerator (RIA) and similar facilities become operational.  These accelerators will allow studies of the GT distributions in unstable nuclei by charge-changing reactions like (d,$^2$He) or ($^{12}$C,$^{12}$N) in inverse kinematics.  Such experiments, as well as high-precision beta-decay measurements will either verify or severely constrain the nuclear models.  They will also determine energies, quantum numbers, and low-lying transition strengths at low excitation energies.  For example, the shell model calculations predict that, with increasing neutron excess, the centroid of the GT distribution moves to lower excitation energies in the daughter nucleus \\cite[see Fig.~8 in]{LM00}.  It would be very helpful to quantify this prediction experimentally. Such data could appreciably improve our understanding of proton-neutron correlations and of the evolution of the $f_{7/2}-f_{5/2}$ spin-orbit splitting in the $pf$-shell.  There are also currently no experimental data regarding the GT$_+$ distribution for odd-odd nuclei in the $pf$-shell.  This is particularly unfortunate given the key role played by capture on $^{60}$Co in the FFN rates.  The shell model electron capture rates for odd-odd nuclei, however, are significantly -- often by more than one order of magnitude at the relevant temperature and density -- lower than the FFN estimates and this is largely responsible for the difference between the FFN and LMP presupernova models.  This is because the shell model places the centroid of the GT strength distribution for odd-odd nuclei at noticeably higher excitation energies in the daughter nucleus."
    },
    "0011/astro-ph0011229_arXiv.txt": {
        "abstract": "Super--sonic turbulence fragments molecular clouds (MC) into a very complex density field with density contrasts of several orders of magnitude. A fraction of the gas is locked into dense and gravitationally bound cores, which collapse as proto--stars. This process can be studied with numerical simulations of super--sonic self--gravitating turbulence.  In this work, we use numerical simulations of magneto--hydrodynamic (MHD), super--sonic, super--Alfv\\'{e}nic and self--gravitating turbulence to compute the mass distribution of collapsing proto--stellar cores, which are selected as local density maxima. We find that the mass distribution of collapsing cores is consistent with the stellar initial mass function (IMF), suggesting that super--sonic turbulence may be responsible for the generation of the IMF.  To support this conclusion we also show that the physical properties of the numerically selected cores are in agreement with the properties of observed NH$_3$ cores and that their magnetic field strength is consistent with Zeeman splitting measurements.  In turbulent MCs, star formation occurs via the gravitational collapse of super--critical cores, formed by the turbulent flow, sub--critical cores being irrelevant for the process of star formation. ",
        "introduction": "Despite early observational evidence of the highly super--sonic nature of molecular cloud (MC) turbulence (Zuckerman \\& Palmer 1974), there have been attempts to simplify the problem of star formation by: i) assuming that super--sonic turbulence behaves like sub--sonic laboratory turbulence; or by ii) focusing on the process of the formation of one single star, neglecting its larger scale environment. The first type of simplification has allowed theoreticians to consider super--sonic turbulence as a source of internal pressure, as suggested by Chandrasekhar (1958) for sub--sonic turbulence (e.g. Bonazzola et al. 1987, 1992). This simplification puts aside the most important characteristic of super--sonic turbulence, that is, the generation of a very fragmented density field by a complex system of turbulent shocks. The second type of simplification has allowed for the generation of analytical and numerical models for the collapse of proto--stars, based on quasi statically evolving proto--stellar cores (e.g. Shu, Adams \\& Lizano 1987), which are hard to reconcile with the turbulent medium from which they are formed. A complementary way to model the dynamics of MCs avoiding the problem of super--sonic turbulence has been the assumption that the observed line width of molecular transitions is not due to turbulence, but to Alfv\\'{e}n waves, that is, oscillations of a rather strong magnetic field (Arons \\& Max 1975; Zweibel \\& Josafatsson 1983; Elmegreen 1985; Falgarone \\& Puget 1986). However, even in this scenario, random motions and super--sonic shocks are generated by super--sonic compressions along magnetic field lines, and so the process of turbulent fragmentation is unavoidable, irrespective of the magnetic field strength.  Recent numerical studies of super--sonic magneto--hydrodynamic (MHD) turbulence have shed new light on the physics of turbulence. The most important results are: i) Super--sonic turbulence decays in approximately one dynamical time, independent of the magnetic field strength (Padoan \\& Nordlund 1997, 1999; MacLow et al. 1998; Stone, Ostriker \\& Gammie 1998; MacLow 1999); ii) The probability distribution function of gas density in isothermal turbulence is well approximated by a Log--Normal distribution, whose standard deviation is a function of the rms Mach number of the flow (Vazquez--Semadeni 1994; Padoan 1995; Padoan, Jones \\& Nordlund 1997; Scalo et al. 1998; Nordlund \\& Padoan 1999; Ostriker, Gammie \\& Stone 1999); iii) Super--sonic isothermal turbulence generates a complex system of shocks that fragment the gas very efficiently into high density sheets, filaments, and cores; iv) Super--Alfv\\'{e}nic turbulence provides a good description of the dynamics of MCs, and an explanation for the origin of dense cores with magnetic field strength consistent with Zeeman splitting measurements. (Padoan \\& Nordlund 1997, 1999).  We call ``turbulent fragmentation'' the process leading to the generation of high density structures (and low density ``voids'') by turbulent shocks. In previous works we investigated the possibility that MCs are fragmented primarily by random super--sonic motions, assuming that self--gravity is important only in the densest regions, so we used numerical simulations of turbulent flows where the effects of self--gravity were intentionally disregarded (Padoan, Jones \\& Nordlund 1997; Padoan et al. 1998, 1999; Padoan \\& Nordlund 1997, 1999; Padoan, Zweibel \\& Nordlund 2000; Padoan et al. 2000). We found that many observational properties of MCs are indeed consistent with the properties of turbulent flows, which supports the suggestion that gravity is important only in the densest regions formed by turbulent shocks.  In this work we include the description of self--gravity, in simulations of three dimensional super--sonic MHD turbulent flows (\\S~2). Self--gravity is here required because we are interested in selecting gravitationally bound (collapsing) protostellar cores, formed by the process of turbulent fragmentation. The mass distribution of proto--stellar cores, selected numerically as local maxima in the gas density, is computed in \\S~3. The value of this result is then confirmed by the comparison between the numerically selected cores and the observed NH$_3$ cores, in \\S~4. The simulation also allows us to predict core properties, which are difficult to measure observationally and are therefore not well known, such as the magnetic field strength (\\S~5). In \\S~6 we discuss the results of this work, and in \\S~7 we list our main conclusions. ",
        "conclusions": "We have computed physical properties of collapsing cores, in a simulation of super--sonic, super--Alfv\\'{e}nic, self--gravitating MHD turbulence. The main results are: i) The mass distribution of collapsing MHD cores is consistent with the stellar IMF; ii) The size, density, line width and velocity gradient of MHD and NH$_3$ cores are very similar; iii) MHD cores and observed NH$_3$ cores show the same correlation between line width and size; iv) Proto--stellar cores with significant magnetic field strength, even in excess of 100~$\\mu$G, are formed in a flow with a low average field strength, $\\langle B\\rangle\\approx 4.5$~$\\mu$G; v) MHD cores show the same relation between magnetic field strength and column density as estimated with Zeeman splitting measurements.  The general conclusion is that turbulent fragmentation is essential to understand the generation of the stellar IMF. Furthermore, star formation, independent of the stellar mass, occurs via the collapse of super--critical cores formed inside super--sonic turbulent MCs, sub--critical cores being irrelevant for the process of star formation. This point of view conflicts with many years of literature on star formation, but it is very promising. It has become accessible to investigation only in the past few years, because of important advances in the field of numerical fluid--dynamics, due to the increased power of super--computers and three--dimensional visualization software."
    },
    "0011/astro-ph0011535_arXiv.txt": {
        "abstract": "AX J2049.6+2939 is a compact X-ray source in the vicinity of the southern blow-up region of the Cygnus Loop supernova remnant (Miyata et al. 1998a). This source was the brightest X-ray source inside the Cygnus Loop observed during the ASCA survey project. The X-ray spectrum was well fitted by a power-law function with a photon index of $-2.1 \\pm 0.1$. Short-term timing analysis was performed and no coherent pulsation was found. Follow-up observations with ASCA have revealed a large variation in X-ray intensity by a factor of $\\simeq$ 50, whereas the spectral shape did not change within the statistical uncertainties. In the second ASCA observation, we found another X-ray source, AX J2050.0+2941, at the north east of AX J2049.6+2939. During the three ASCA observations, the X-ray intensity of AX J2050.0+2941 varied by a factor of $\\simeq$4. No coherent pulsations could be found for AX J2050.0+2941.   We have performed optical photometric and spectroscopic observations in the vicinity of AX J2049.6+2939 at the Kitt Peak National Observatory (KPNO). As a result, all objects brighter than $B$-band magnitude of 22 in the error box can be identified with normal stars. Combined with the X-ray results and the fact that there are no radio counterparts, AX J2049.6+2939 is not likely to be either an ordinary rotation-powered pulsar or an AGN. The nature of AX J2049.6+2939 is still unclear and further observations over a wide energy band are strongly required.  As to AX J2050.0+2941, the long-term X-ray variability and the radio counterpart suggests that it is an AGN. ",
        "introduction": "Supernovae (SN) with a massive progenitor star leave behind a stellar remnant. A massive star with a mass ($M$) of 10 \\Msun $< M < 40$ \\Msun\\ leaves a neutron star while that with a M $>$ 40 \\Msun\\ leaves a blackhole.  The Cygnus Loop is a well evolved supernova remnant (SNR) and has been widely observed from radio to X-ray wavelengths. Miyata et al. (1998b) investigated the central portion of the Cygnus Loop and concluded that there were Si, S, and Fe rich plasmas in the center of the Loop.  Comparing their abundance ratios with the theoretical calculations of nucleosynthesis by a SN explosion, they estimated the progenitor mass of $\\sim$25 \\Msun.  Furthermore, Miyata et al. (1999) reported a progenitor mass of $\\sim$15 \\Msun\\ based on the radial distribution of heavy elements from the north-eastern region toward the central portion. In both cases, we expect that a neutron star must have been left behind the SN which produced the Cygnus Loop. As mentioned in Miyata et al. (1998a; hereafter paper I), there have been several observations to search for a neutron star associated with the Cygnus Loop.  So far, no compact source has been established.  The ASCA Observatory is a powerful tool to discover embedded neutron stars because of its high detection efficiency at hard X-ray region and has discovered more than a dozen neutron stars. We have reported the discovery of a compact X-ray source, AX J2049.6+2939, during the ASCA survey project of the Cygnus Loop (Paper I). Its X-ray properties suggest that it is a neutron star or an active galactic nucleus (AGN). If it is a neutron star, it is definite that the Cygnus Loop must have originated from a massive progenitor star.  In this paper, we performed optical photometric and spectroscopic observations to investigate whether or not it is an AGN. We also performed the timing analysis for follow-up ASCA observations and RXTE observations to investigate the possibility of a neutron star which might associate with the Cygnus Loop. ",
        "conclusions": "We performed wide-band spectroscopic observations in the vicinity of AX J2049.6+2939. Based on the optical photometric and spectroscopic observations, we can rule out the possibility of an AGN. AX J2049.6+2939 shows a large X-ray variability over 19 years, which excludes the possibility of an ordinary rotation-powered pulsar.  Thus, the ADAF model would be a plausible one, although no definite evidence is obtained.  As calculated by Narayan, Barret, \\& McClintock (1997) and Manmoto, Mineshige, \\& Kusunose (1997), the spectrum generated by ADAF shows the synchrotron peak which comprises the Rayleigh-Jeans slope and the optically thin synchrotron emission in the radio regime. The Comptonization of the synchrotron photons leads to one or two peaks from IR to UV regimes. In the X-ray and $\\gamma$ ray regimes, the bremsstrahlung emission plus photons suffering multiple Compton scattering and emission from saturated Comptonized photons are dominant. Therefore, further simultaneous observations over a wide energy band are essential to investigate its nature.  We have found another hard X-ray emitting compact source, AX J2050.0+2941, which is in a direction north-eastern from AX J2049.6+2939.  We also found radio (4850MHz) emission in the vicinity of AX J2050.0+2941.  Due to its large error circle, we found more than forty objects for the optical candidate of AX J2050.0+2941. We therefore need further X-ray observations in the vicinity of AX J2050.0+2941 with Chandra or XMM-Newton in order to determine its position more accurately.  It would be essential to investigate the relationship with the Cygnus Loop.  \\vspace{1pc}\\par We thank Profs. S. Kitamoto, F. Nagase, and S. Mineshige for invaluable comments and suggestions.  We appreciate Mr. C. Baluta for careful reading of manuscript. We thank {\\sl ASCA\\_ANL} and {\\sl DISPLAY 45} developing team members, especially Dr. Ishisaki.  Part of this research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. This research is partially supported by ACT-JST Program, Japan Science and Technology Corporation.  \\clearpage"
    },
    "0011/astro-ph0011473_arXiv.txt": {
        "abstract": "According to Pei, Fall \\& Hauser (1999), the global metallicity evolution of the Universe can be represented by the ratio of the total metal content to the total gas content measured in Damped Lyman--$\\alpha$ (DLA) systems (the ``column density weighted metallicity'' \\`a la Pettini). To minimize dust obscuration effects, a DLA sample with negligible dust content is considered, namely, 50 DLAs with $\\log N_{HI} < 20.8$. The global metallicity found shows clear evidence of redshift evolution that goes from $\\sim1/30$ solar at $z\\sim4.1$ to solar at $z\\sim0.4$. More generally, DLAs with measured heavy elements probe the ISM of high redshift galaxies. The whole sample collected from the literature contains \\nnhi~DLAs. The metallicity is calculated adopting for the dust correction the most general method used so far, based on models of the ISM dust depletions in the Galaxy.  The intrinsic metallicity evolution of DLA galaxies is $d\\log Z_{DLA}/dz \\propto -0.33 \\pm 0.06$. ",
        "introduction": "A Damped Lyman--$\\alpha$ (DLA)  is often referred to as a QSO absorption line system with HI column density larger than $2\\times10^{20}$ atoms cm$^{-2}$ associated with clouds in protogalactic disks (Wolfe et al.~1986). In a more modern fashion, DLAs are mostly neutral gas clouds (HI column density larger than $\\sim10^{19}$ atoms cm$^{-2}$) associated with the interstellar medium (ISM) of a not defined class of galaxies detected from redshift $z=0$ up to $z=4.6$. Among others, examples of the variety of morphologies of the emitting counterpart found nearby DLAs are given by Le Brun et al.~(1997) and Rao \\& Turnshek (1998).  The large investment of telescope time dedicated to QSO surveys (Fan et al.~1999) and specifically to DLA surveys (Rao \\& Turnshek 1999) has brought the discovery of numerous DLAs. We know \\ncon~ objects (Fig.~1) and \\ncan~ candidates for a total of \\ntot~DLAs in the redshift range $0.0 < z < 4.6$. Among these, \\nnhi~are DLAs for which the column density of HI and one or more of the following ions have been measured: FeII, SiII, NiII, MnII, CrII, ZnII.  These ions are the dominant contributors to the abundances of the corresponding elements in HI clouds with high column densities, because they all have ionization potentials below 13.6~eV.  For this reason it is easy and straightforward to determine the relative heavy element abundance, just using the standard relation [X/Y] = $\\log (N_X/N_Y)_{DLA} - \\log (X/Y)_\\odot$.  An important contribution to the richness of this sample has been provided by Prochaska \\& Wolfe (1999), Pettini et al.~(1997), and Lu et al.~(1996). The heavy element enrichment of DLAs has been studied recently in detail by Pettini et al.~(1999) and Vladilo et al. (2000) who considered the ZnII absorption as a tracer of metallicity, and by Prochaska \\& Wolfe (2000), who preferred FeII measurements instead. Here we try to show that by combining this information together to build a larger database (\\nnhi~DLAs at $0.0<z<4.4$) and making some assumptions about biases due to dust obscuration, we can have good evidence that the metallicity of the Universe as given by DLAs evolves with cosmic time.  \\begin{figure} \\plotone{fig1.ps} \\caption{Neutral hydrogen column density for a sample of \\ncon~DLAs in the redshift range $0.0<z<4.6$.  Dots are DLAs with at least one measured heavy element, triangles are the remaining DLAs. Those points without error bars are DLAs with no HI uncertainty measured.} \\end{figure} ",
        "conclusions": "It is not clear yet how well the mean HI weighted metallicity obtained from current observations can represent the global metal enrichment of the Universe because different observational biases might affect the results.  As shown in Fig.~8, low and high $N_{HI}$ DLAs are distributed in a different way (the latter are concentrated at the center of the redshift interval while the former is much more scattered). Moreover, the two correlations are consistent with each other at the $\\sim2\\sigma$ level.  The lack of high $N_{HI}$ with high metallicity, theoretically first considered by Pei and Fall (1995) and then found by Boiss\\'e et al. (1998), might correspond to the lack of high $N_{HI}$ DLAs at $z<1.5$. If there were a metallicity evolution also for high $N_{HI}$ DLAs, dust obscuration could be too large at $z<1.5$ to make DLAs detectable (a DLA with $\\log N_{HI}=21.0$ and metallicity $\\log Z_{DLA}/Z_\\odot=-0.5$ would result in an extinction of about 1 magnitude at Ly$\\alpha$). Moreover, this is where UV observations, required to detect Ly$\\alpha$ absorption, are very sensitive to any appreciable dust opacity.  In general, before using DLAs to explore the metallicity evolution in the Universe, one should understand if and how DLA galaxies represent the total galaxy population at different epochs. Obviously, this is not a trivial task because many factors have to be taken into account some of which are very poorly known. When trying to match Lyman break galaxies to DLA galaxies, one should consider that the former are star forming galaxies for which the metallicity is usually inferred (and not measured) from the stellar and gaseous emissivity, while the latter are interstellar clouds with large HI column densities in galaxies of undefined nature for which the metallicity is measured from the absorption lines. On the other hand, when comparing DLAs with the intergalactic medium, one has to consider that high density regions like DLAs have very likely a different metallicity history than the mean value in the Universe and/or in the diffuse intergalactic medium. Cen \\& Ostriker (1999) have pointed out that higher density regions have higher star formation rates and after a fast heavy--element enrichment, the metallicity evolves very slowly compared to lower density regions, due to a suppression of the star formation.  We like to stress that it is the combination of a large sample of DLAs (\\nnhi~objects), a wide redshift baseline ($0.0<z<4.4$) and an accurate dust correction applied to all heavy elements measured, that has allowed us to demonstrate that the metallicity of DLAs evolves with redshift. Had we limited ourselves to using Zn only (detected in 31 DLAs) with no dust correction, we would have not detected the same effect, even if ZnII lines are known not to suffer from strong saturation and dust depletion. This is because even small corrections are important when trying to detect elusive effects. In addition, Zn is much less abundant than other elements such as Fe and Si and, therefore, measuring its absorption lines is not always an easy task. The combined use of as many elements as possible allows one to much more efficiently track the metallicity evolution down to gas clouds that are very metal poor."
    },
    "0011/astro-ph0011190_arXiv.txt": {
        "abstract": "We present a strategy to identify several thousand stars that are astrometrically stable at the micro-arcsecond level for use in the \\itsim\\ (Space Interferometry Mission) astrometric grid. The requirements on the grid stars make this a rather challenging task. Taking a variety of considerations into account we argue for K~giants as the best type of stars for the grid, mainly because they can be located at much larger distances than any other type of star due to their intrinsic brightness. We show that it is possible to identify suitable candidate grid K~giants from existing astrometric catalogs. However, double stars have to be eliminated from these candidate grid samples, since they generally produce much larger astrometric jitter than tolerable for the grid. The most efficient way to achieve this is probably by means of a radial velocity survey. To demonstrate the feasibility of this approach, we repeatedly measured the radial velocities for a pre-selected sample of 86~nearby \\hipparcos\\ K~giants with precisions of 5--8\\,\\ms. The distribution of the intrinsic radial velocity variations for the bona-fide single K~giants shows a maximum around 20\\,\\ms, which is small enough not to severely affect the identification of stellar companions around other K~giants. We use the results of our observations as input parameters for Monte-Carlo simulations on the possible design of a radial velocity survey of all grid stars. Our favored scenario would result in a grid which consists to 68\\% of true single stars and to 32\\% of double or multiple stars with periods mostly larger than 200~years, but only 3.6\\% of all grid stars would display astrometric jitter larger than 1\\,\\muas. This contamination level is probably tolerable. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011159_arXiv.txt": {
        "abstract": "{ We present the temporal analysis of X-ray observations of the radio-loud Narrow-Line Seyfert 1 galaxy (NLS1) \\object{PKS~0558--504} obtained during the {\\it XMM-Newton} Calibration and Performance Verification (Cal/PV) phase. The long term light curve is characterized by persistent variability with a clear tendency for the X-ray continuum to harden when the count rate increases. Another strong correlation on long time scales has been found between the variability in the hard band and the total flux. On shorter time scales the most relevant result is the presence of smooth modulations, with characteristic time of $\\sim$ 2 hours observed in each individual observation. The short term spectral variability turns out to be rather complex but can be described by a well defined pattern in the hardness ratio--count rate plane. ",
        "introduction": "Active Galactic Nuclei (AGN) are variable in every observable wave band. The X-ray flux exhibits variability on shorter time scales than any other energy band, indicating that the emission occurs in the innermost regions of the central engine. Therefore, a study of the X-ray variability is a powerful tool providing upper limits on the sizes of the emitting regions and the masses of central black holes and it allows one to probe the extreme physical processes operating in the inner parts of the accretion flow close to the event horizon. Although X-ray variability has been observed in AGN for more than two decades, its origin is still poorly understood. Narrow-Line Seyfert 1 galaxies (NLS1) often display rapid, large amplitude X-ray variability as well as extreme long term changes (Forster \\& Halpern \\cite{forst}, Boller \\cite{boll1}, Brandt \\cite{brand}), and therefore they represent the ideal class for an X-ray temporal analysis. \\begin{table} \\caption{EPIC PN individual observations.} \\begin{center} \\begin{tabular}{llll} \\hline \\hline \\noalign{\\smallskip} Observation date & Duration& Mode$^a$ & Filter$^b$\\\\ \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} 00/2/7 UT11:15-13:19 & 7.2 ks & FF & M\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00/2/7 UT14:28-16:02 & 5.4 ks & LW & M\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00/2/10-11 UT23:26-3:23 & 14.0 ks & FF & T2\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00/3/2 UT18.15-21:42 & 12.4 ks & FF & M\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00/5/24 UT17:21-21:07 & 13.6 ks & SW & M\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} $^a$ Full Frame (FF), Large Window (LW), Small Window (SW).\\\\ $^b$ Medium (M), Thin 2 (T2).\\\\ \\end{table} \\object{PKS~0558--504} ($z=0.137,~m_{\\rm B}=14.97$) is one of the few radio-loud NLS1. In X-rays it is characterized by a steep spectrum [the photon index $\\Gamma$ ranges between 2.2 and 3.1, obtained from different instruments and energy bands (a detailed spectral analysis of {\\it XMM-Newton} observations has been performed in a separate paper, O'Brien \\cite{obri})], high luminosity [$(2-5)\\times 10^{45}{~\\rm erg~s^{-1}}$, with $H_0=70 {~\\rm km~s^{-1}~Mpc^{-1}},~ q_0=0.5$], and strong variability. A Ginga observation (Remillard \\cite{remi}) showed an increase of the X-ray flux by $67\\%$ in 3 minutes, implying the presence of a relativistic flare. Recent ROSAT High Resolution Imager (HRI) observations (Gliozzi \\cite{glioz}) confirmed the presence of strong and persistent X-ray variability, which suggests the presence of a rotating black hole, and ruled out external contributions to the high luminosity and variability from other nearby sources.  In this paper we report the results of four {\\it XMM-Newton} observations of \\object{PKS~0558--504} taken with  the European Photon Imaging Camera (EPIC) PN, which is the best suited instrument on board the ESA satellite for timing analysis purposes, due to its high time resolution. In Sect.2 we describe the observations and data reduction. The long and short term X-ray variability are discussed in Sect.3 and Sect.4, respectively. Sect.5 contains the main conclusions and a summary. ",
        "conclusions": "We have presented the results of the EPIC PN observations of the NLS1 \\object{PKS~0558--504} taken during the Cal/PV phase. We found: \\begin{itemize}  \\item During the XMM pointings \\object{PKS~0558--504} showed a moderate but persistent variability both on long (months) and short (minutes and hours) time scales.  \\item Strong correlations between hardness ratio and total flux have been found. When the mean flux is above a certain threshold, the spectrum becomes harder as the count rate increases. When the flux is below a certain level, the spectral behaviour is in the opposite sense. At intermediate count rates no significant spectral variability is observed.  \\item A further strong correlation between the broad band flux and the variability (expressed in terms of excess variance) in the hard energy band has been discovered: the contribution to the total variability of the hard band increases and becomes dominant as the mean count rate raises.  \\item In all the observations the variability is characterized by smooth similar modulations, with typical time scale $\\sim 2.2$ hr, determined with the structure function analysis.  \\item The most extreme temporal variation of the count rate implies a radiative efficiency slightly larger than the theoretical maximum for accretion onto a Schwarzschild black hole. It is important to point out that this is a strictly lower limit, since during our analysis only single events (which represent $\\sim 70\\%$ of the total X-ray photons collected) have been used. \\end{itemize}  A detailed model of \\object{PKS~0558--504} which accounts for all timing and spectral features discovered is beyond the scope of this paper. Therefore, we limit ourselves to some general considerations. The most intriguing result is the tentative detection of a common characteristic time scale of $\\sim $ 2 hours. However a confirmation of this characteristic time scale has to await for longer observations, to sample a higher number of cycles. A more detailed analysis to assess its significance and robustness (e.g. as in Boller \\cite{boll2}) is deferred to a follow-up publication. This time scale might be associated to periodic phenomena occurring in the accretion disk, whose contribution seems to be dominant also on the basis of the spectral analysis of EPIC MOS data (O'Brien \\cite{obri}). Two possible explanations which need to be quantitatively tested  are: {\\it i}) the presence of a hot spot orbiting the black hole and {\\it ii}) the precession of the inner part of the accretion disk caused by misalignment of black hole and disk rotational axes. An alternative tempting hypothesis involves the presence of a jet: as \\object{PKS~0558--504} is a radio-loud object, beamed emission from a jet could partly contribute to the brightness and variability in X-rays. Radio-loud quasars are known to have flatter X-ray spectral indices than radio-quiet ones, and they are also known to be brighter X-ray sources (e.g. Yuan \\cite{yuan}). A widely accepted explanation for this fact is that the X-ray emission includes an additional harder component associated with the radio jet. In this scenario a naive interpretation of the temporal and the spectral variability associated with changes of the broad band flux, involves the presence of either a precessing jet or moving knots with helical trajectories in a relativistic magnetized jet (Camenzind \\& Krockenberger \\cite{came}). At this point it is worth noticing that also in one other of the few radio-loud NLS1, RX J0134.2--4258, the spectral variability (characterized, as in our case, by a hardening of the spectrum during the flux increase) was tentatively interpreted as  due to the jet emission (Grupe \\cite{grupe}).  The exciting results from the temporal analysis of a few short exposures demonstrate the extraordinary capabilities of XMM, in particular in light of the long observations possible in virtue of the highly eccentric orbit."
    },
    "0011/astro-ph0011409_arXiv.txt": {
        "abstract": "Embedded massive starclusters have recently been identified in several nearby galaxies by means of the radio-wave thermal bremsstrahlung emission from their surrounding \\HII\\ regions.   Energy requirements imply that these optically-obscured starclusters contain 500-1000 O-type stars, making them similar to the ``super starclusters'' observed in many dwarf starbursts and mergers.  Based on their high free-free optical depth and visual extinctions of $A_V\\gg10$ mag., these massive ``ultra-dense'' \\HII\\ regions  (UDHIIs) are distinct signatures of the youngest, most compact super starclusters.  UDHII regions may represent the earliest stages of globular cluster formation.    We review the properties of presently-known UDHIIs, and we outline a pictoral evolutionary taxonomy for massive cluster formation which is analogous to the more familiar evolutionary sequence for individual stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011123_arXiv.txt": {
        "abstract": "We give for a plasma with a history of several expansion and contraction phases an analytical model of the evolution of a contained relativistic electron population under synchrotron, inverse Compton and adiabatic energy losses or gains. This is applied to different scenarios for evolution of radio plasma inside the cocoons of radio galaxies, after the activity of the central engine has ceased. It is demonstrated that fossil radio plasma with an age of even up to 2 Gyr can be revived by compression in a shock wave of large-scale structure formation, caused during the merging events of galaxy clusters, or by the accretion onto galaxy clusters.  We argue, that this is a highly plausible explanation for the observed cluster radio relics, which are the regions of diffuse radio emission found in clusters of galaxies, without any likely parent radio galaxy seen nearby. An implication of this model is the existence of a population of diffuse, ultra-steep spectrum, very low frequency radio sources located inside and possibly outside of clusters of galaxies, tracing the revival of aged fossil radio plasma by the shock waves associated with large-scale structure formation. ",
        "introduction": "The radio cocoons of radio galaxies become rapidly undetectable after the central engine of the active galactic nucleus ceases to inject fresh radio plasma, or even earlier \\cite[and references therein]{2000MNRAS.315..381K}.  Radio plasma consists of relativistic electrons and magnetic fields, which are responsible for the synchrotron radio emission, and possibly relativistic protons and/or a non-relativistic thermal gas component.  Although undetectable, aged radio plasma should still be an important component of the inter-galactic medium (IGM). We investigate the possibility of reviving patches of such fossil radio plasma (also called {\\it radio ghosts}, En{\\ss}lin \\cite*{Ringberg99}) by compression in a shock wave produced in the IGM by the flows of cosmological large-scale structure formation. Such shock waves re-energize the electron population in the fossil radio cocoons, which can lead to observable synchrotron emission. This is probably the mechanism responsible for the, so called, {\\it cluster radio relics}, which are patches of diffuse, sometimes polarized radio emission typically found at peripheral locations in clusters of galaxies. These cluster radio relics can not be simply relic radio galaxies, as their name suggests.  The spectral ages of the electron population are {usually} too short to admit even the nearest galaxy to have been the parent radio galaxy, which has moved to its present location with a velocity typical for cluster galaxies. Therefore, a recent enhancement of the nonthermal radio output of the cluster relic sources is mandatory.  For reviews on diffuse cluster radio emission (cluster radio relic and halos, which are believed to be distinct phenomena), see Jaffe \\cite*{1992cscg.conf..109J}, Feretti \\& Giovannini \\cite*{1996IAUS..175..333F}, En{\\ss}lin et al. \\cite*{1998AA...332..395E}, Feretti \\cite*{Feretti.Pune99}, En{\\ss}lin \\cite*{Pune99}, and Giovannini et al. \\cite*{1999NewA....4..141G}.  The connection between the presence of a shock wave and the appearance of the cluster radio relic phenomena was assumed to be due to Fermi-I shock acceleration of electrons, in En{\\ss}lin et al. \\cite*{1998AA...332..395E} and Roettiger et al. \\cite*{1999ApJ...518..603R}. While this process might work within the normal IGM, several arguments favor the possibility that indeed old fossil radio plasma is revived in the case of cluster relic sources: \\begin{itemize} \\item Cluster radio relics are extremely rare, whereas shock waves should be very common within clusters of galaxies. The {dual} requirement of a shock wave and fossil radio plasma, {for producing} a cluster radio relic, {would} be an attractive explanation for the rareness of the relics. \\item Fossil radio plasma with existing relativistic electron population and fairly strong magnetic field appears to have ideal properties to be brightened up during the shock's passage. \\item The cluster radio relic 1253+375 near the Coma cluster of galaxies appears to be fed with radio plasma by the nearby galaxy NGC 4789 (see Fig \\ref{fig:ngc4789} and scenario C in Sec. \\ref{sec:scenarios}). \\end{itemize}   But, if indeed the fossil radio plasma and not the normal IGM were to become radio luminous at a shock wave, the expected very high sound velocity of that relativistic plasma should forbid the shock in the ambient medium to penetrate into the radio plasma. Thus, shock acceleration is not expected to occur there. Instead, the fossil radio plasma would get adiabatically compressed, and the energy gain of the electrons is expected to be mainly due to adiabatic heating. It is the purpose of this work to demonstrate that this process is sufficient to account for the cluster radio relics.  {Sec. \\ref{sec:shocks} and \\ref{sec:fossil} discuss the expected properties of the components of this model: shock waves and fossil radio plasma.}  In Sec. \\ref{sec:formalism} we give a mathematical formalism to describe the evolution of the radiation spectrum during the history of a radio cocoon. The different phases of this history are described in Sec. \\ref{sec:phases} and applied to three different scenarios in Sec. \\ref{sec:scenarios}. The discussion of the merits and potential deficiencies of this model can be found in Sec. \\ref{sec:discussion}.  \\begin{figure}[t] \\begin{center} \\psfig{figure=flash.fig1.eps,width=0.5 \\textwidth,angle=0} \\end{center} \\caption[]{\\label{fig:ngc4789} Radio map of the cluster radio relic 1253+275 in the Coma cluster from Giovannini et al. \\cite*{1991A&A...252..528G} and its probable source of radio plasma: the narrow angle tailed radio galaxy NGC 4789. The radio tails of NGC 4789 indicate the direction from which this galaxy came: the cluster center which is located outside this map. A radio polarization of 27 \\% was detected from the radio relic \\cite{1991A&A...252..528G}, but not from the tails of NGC 4789 feeding the relic with radio plasma. This is a clear signature of the alignment of magnetic fields enhanced in the shock compression. The polarization vectors are consistent with a shock wave oriented parallel to the main axis of the radio relic \\cite{1998AA...332..395E}.} \\end{figure} ",
        "conclusions": "}  We have argued here that adiabatic compression in cluster shocks can revive fossil radio plasma to radio detection, even up to 2 Gyr after its release from the parent radio galaxy. The computed radio spectra provide a good match to the observed spectra of cluster radio relics. Hence, this is a promising model for the regions of diffuse radio emission found in clusters of galaxies. Below, we summarize some merits and potential shortcoming of this model:\\\\  \\noindent {\\bf Pros:} \\begin{itemize} \\item The observed connection of cluster radio relics to shock waves arises naturally in this model. \\item The presence of the radio galaxy NGC 4789 and the morphological connection of its radio tails to the relic 1253+275 seems to be the {\\it smoking gun} of cluster radio relic formation by compression of fossil radio plasma. \\item The expected very high sound speed within radio plasma should virtually forbid shock waves of the ambient medium to penetrate into radio cocoons. This renders adiabatic compression as a more plausible means of reviving the fossil plasma. \\item Cluster radio relics are rare. Not only a shock wave and fossil radio plasma have to be present in this model, in order to produce a cluster radio relic, but also the fossil radio plasma can not grow too old (of the order of 0.1 Gyr in the center of clusters and 1 Gyr at their boundaries). Otherwise the high energy end of the electron population would suffer too severe a depletion for the adiabatic gains during the compression to be able to shift them to radio emitting energies. Since fossil radio plasma should be very common \\cite{ensslin2000a}, this helps to explain the observed rarity of cluster radio relics. \\item The observed tendency of relics to appear at peripheral positions in clusters follows in this model from the much shorter lifetimes of the radio emission and, consequently, much shorter ages which the revive-able fossil radio cocoons can have in the center of clusters, owing to the stronger synchrotron losses occurring there. \\end{itemize}  \\noindent {\\bf Contras:} \\begin{itemize} \\item If the radio plasma has a high mass load, due to undetectable cool gas, it should get shocked, ensuing Fermi-I shock acceleration {to produce the observed radio emitting electron population \\cite{1998AA...332..395E}}.  Still, the adiabatic gains of the relativistic particles are inescapable, although restricted in this case due to the limited compression factor of a shock wave. \\item {In the case when the fossil radio plasma is supported by a very hot, non-relativistic gas component, both adiabatic and shock acceleration are ineffective. The first due to the lower compressibility of a non-relativistic equation of state, while the second due to the inability of environmental shock waves to penetrate into the plasma. The fact that we do observe cluster radio relics tells us that this situation does not necessarily arise (or, alternatively, cluster radio relics are not related to radio ghosts).} \\item The giant cluster radio relic 2006-56 in Abell 3667 shows a rim of flat radio emission (spectral index around 0.5) at its sharp outer edge \\cite{1997MNRAS.290..577R}. This could be understood in a Fermi-I shock acceleration model due to the different ages of the electrons found at different distances from the energizing shock wave \\cite{1998AA...332..395E}. There are two possibilities to explain the observation of a flat spectrum rim within the adiabatic compression model: (a) there is a high mass load in the radio plasma, allowing the shock wave to successively compress the plasma, which should also produce a spatially different electron population, and (b) the apparent rim of a flat spectral index is an artifact arising from an unmatched coverage of the Fourier space in the radio interferometric observations at the two frequencies used. \\end{itemize} We conclude that the model presented here provides a fairly natural and promising explanation for the phenomenon of cluster radio relics. This model predicts the existence of a population of diffuse, ultra-steep, very low frequency radio sources inside and possibly also outside of clusters of galaxies, due to the age dependence of the upper frequency cutoff of radio emission arising from the revivable fossil radio plasma. The ongoing rapid improvements in the sensitivities at low radio frequencies (Giant Meterwave Radio Telescope: (GMRT), Low Frequency Array (LOFAR)) may therefore open a new window on cosmological structure formation by detecting shock waves marked by such relic radio sources."
    },
    "0011/astro-ph0011065_arXiv.txt": {
        "abstract": "We present results of  three dimensional simulations of the uppermost part of the sun, at 3 stages of its evolution. Each  model includes physically realistic radiative-hydrodynamics (the Eddington approximation is used in the optically thin region), varying  opacities and a realistic equation of state (full treatment of the ionization of H and He). In each evolution model, we investigate a domain, which starts at the top of the photosphere and ends just inside the convection zone (about 2400 km in the sun model). This includes all of the super-adiabatic layer (SAL). Due to the different positions of the three models in the $log (g) $ vs $log T_{eff}$ plane, the more evolved models have lower density atmospheres. The reduction in density causes the amount of overshoot into the radiation layer, to be greater in the more evolved models. ",
        "introduction": "In most of the convection zone, heat transport can be modelled quite well by using the mixing length theory. Due to vigorous turbulent mixing, the entropy is almost constant. This well mixed layer has an extremely small super-adiabatic temperature gradient ($\\na - \\na_{\\rm ad}  \\approx 10 ^{-8}$) and is practically opaque (the optical depth is greater than $10^4$).  In such a layer,  radiative transport accounts for a very small fraction of the heat flow. This changes further up. Near the top of the convection zone the gas density is very low, so the enthaply flux is not big enough to carry all of the outward flowing heat.  Radiation must carry most of the heat flux. To enable radiation to transport this heat flux, the temperature gradient must increase significantly. This region, in which $\\na -\\na_{ad}$ is of order unity, is known as the super adiabatic layer (SAL).  The SAL is important for the following reasons:  \\begin{itemize} \\item It is the site of maximum amplitude of the p-mode oscillations (particularly the high frequency modes). \\item As most of the entropy change occurs within this layer, it plays an important role in  determining the radius of stellar models (Larson 1974). (Note $dS/dlnP=c_{\\rm p} [\\na - \\na_{\\rm ad}]$.) \\item The amount of overshoot into the enveloping photosphere depends on the  structure of the SAL. \\end{itemize}  In this poster we will describe three simulations of stellar atmospheres each centered around the SAL. We consider the sun at three different stages of its evolution, namely the ZAMS, present sun and subgiant. These are shown as three points in the theoretical HR diagram, figure \\ref{hr}.  \\begin{figure} \\centering \\epsfig{file=hr1.eps, width=8cm} \\caption{Evolutionary track for the sun from the pre-main sequence stage to the giant branch} \\label{hr} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011071_arXiv.txt": {
        "abstract": "We performed $^{12}$CO(1 -- 0), $^{13}$CO(1 -- 0), and HCN(1 -- 0) interferometric observations of the central region (about 450 pc in radius) of M82 with the Nobeyama Millimeter Array, and have successfully imaged a molecular superbubble and spurs. The center of the superbubble is clearly shifted from the nucleus by 140 pc. This position is close to that of the massive black hole (BH) of $\\gtrsim460$ M$_{\\odot}$ and the $2.2\\mu$m secondary peak (a luminous supergiant dominated cluster), which strongly suggests that these objects may be related to the formation of the superbubble. Consideration of star formation in the cluster based on the infrared data indicates that (1) energy release from supernovae can account for the kinetic energy of the superbubble, (2) the total mass of stellar-mass BHs available for building-up the massive BH may be much higher than 460 M$_{\\odot}$, and (3) it is possible to form the middle-mass BH of $10^{2}-10^{3}$ M$_{\\odot}$ within the timescale of the superbubble. We suggest that the massive BH was produced and is growing in the intense starburst region. ",
        "introduction": "\\label{intro}  A starburst is a star formation event with a high star formation efficiency. This usually occurs in the galactic nuclear/central regions. Starbursts and successive supernova explosions cause drastic changes in kinematic and physical conditions of interstellar medium, and produce (expanding) superbubbles, chimneys, and/or large-scale outflows. There is observational evidence for these structures in hot and/or atomic gas in many galaxies \\citep[e.g.][]{fab89,ten88}, but there is much less clear evidence in molecular gas. Starburst phenomena in the nuclear regions are also thought to be closely related to active galactic nuclei (AGNs) powered by massive black holes (MBHs; more massive than a stellar-mass BH), such as fueling AGNs and/or causing the growth of MBHs. One question, ``can starbursts produce the seeds of MBHs or middle-mass BHs?,'' is still largely unresolved.  Nearby \\citep[3.25 Mpc;][]{san75} irregular galaxy M82 (NGC 3034) has kilo-parsec scale bipolar outflows which can be seen by optical emission lines \\citep[e.g.][]{sho98} and X-ray \\citep[e.g.][]{bre95} (Fig.~\\ref{bubble}a). This outflow is believed to be made by frequent supernova explosions at the central region in consequence of starburst. Recent X-ray studies have found evidences of an MBH with a mass of $\\gtrsim460 M_{\\odot}$ in the starburst region \\citep{mat99,pta99}. In this region, however, strong dust absorption prevents detailed optical/near-infrared observations, so we performed millimeter-wave observations which would not be affected by the dust. ",
        "conclusions": ""
    },
    "0011/astro-ph0011247_arXiv.txt": {
        "abstract": "The Accretion-Ejection Instability (AEI), which can occur in magnetized disks near equipartition, is a good candidate to explain the low-frequency QPO in black-hole binaries.  Here we present analytical work concerning the behavior of QPO frequency and the emission of Alfv\\'en waves from the disk to the corona. ",
        "introduction": "The AEI is a spiral instability, similar to galatic spirals but driven by magnetic stress rather than self-gravity.  It occurs in the inner region of an accretion disk threaded by a vertical magnetic field of the order of equipartition with the gas pressure.  The spiral extracts energy and angular momentum from the disk, causing accretion, and stores them in a Rossby vortex at its corotation radius (see also the contributions of Tagger and Caunt, these proceedings).  This vortex then leaks energy and angular momentum as Alfv\\'en waves to the corona, where it can power a wind or a jet. For more details see Tagger,M.\\& Pellat,R. (1999). ",
        "conclusions": "The frequency of the AEI, and its variations with the disk inner radius, make it a very good candidate to explain the ``ubiquitous'' QPO. The accretion energy and momentum extracted from the disk are propagated to the corona as Alfv\\'en waves."
    },
    "0011/astro-ph0011137_arXiv.txt": {
        "abstract": "\\begin{sloppypar} We present photometric ISO 60 and 170 \\um\\ measurements, complemented by some IRAS data at 60 \\um , of a sample of 84 nearby main-sequence stars of spectral class A, F, G and K in order to determine the incidence of dust disks around such main-sequence stars.  Fifty stars were detected at 60\\um; 36 of these emit a flux expected from their photosphere while 14 emit significantly more. The excess emission we attribute to a circumstellar disk like the ones around Vega and $\\beta$ Pictoris. Thirty four stars were not detected at all; the expected photospheric flux, however, is so close to the detection limit that the stars cannot have an excess stronger than the photospheric flux density at 60\\,\\um.  Of the stars younger than 400 Myr one in two has a disk; for the older stars this is true for only one in ten. We conclude that most stars arrive on the main sequence surrounded by a disk; this disk then decays in about 400 Myr.  Because (i) the dust particles disappear and must be replenished on a much shorter time scale and (ii) the collision of planetesimals is a good source of new dust, we suggest that the rapid decay of the disks is caused by the destruction and escape of planetesimals. We suggest that the dissipation of the disk is related to the heavy bombardment phase in our Solar System.  Whether \\textit{all} stars arrive on the main sequence surrounded by a disk cannot be established: some very young stars do not have a disk. And not all stars destroy their disk in a similar way: some stars as old as the Sun still have significant disks.  \\end{sloppypar} ",
        "introduction": "In 1983, while using standard stars to calibrate the IRAS photometry, \\citet{auma:84} discovered that Vega ($\\alpha$ Lyr), one of the best calibrated and most used photometric standards in the visual wavelength range, emits much more energy at mid- and far-infrared wavelengths than its photosphere produces.  Because the star is not reddened Aumann et al. proposed that the excess IR radiation is emitted by small, interplanetary-dust particles in a disk rather than in a spherical envelope.  This proposal was confirmed by \\citet{smit:84} who detected a flat source of scattered light around $\\beta$ Pic, one of the other Vega--like stars detected in the IRAS data \\citep{gill:86}, and the one with the strongest excess. The disk around Vega and other main-sequence stars is the remnant of a much stronger disk built up during the formation of the stars. \\citet{auma:84} pointed out that such disks have a lifetime much shorter than the stellar age and therefore need to be rebuilt continuously; collisions between asteroids are a probable source of new dust \\citep{weis:84}.  Except for the somewhat exceptional case of $\\beta$ Pic \\citep{hobb:85} and in spite of several deep searches no trace of any gas has ever been found in the disks around main-sequence stars; see e.g. \\citet{lise:99}.  Since 1984 the search for and the study of remnant disks has made substantial progress by the discovery of numerous ``fatter'' disks around pre-main sequence stars that contain dust and gas; for overviews see \\citet{beck:96}, \\citet{sarg:93}, \\citet{vand:98}. The IRAS data base also contains detections of Vega-like disks around red giant stars that have developed from A and F-type main-sequence stars \\citep{plet:97}.  The discovery of Aumann et al. has prompted deeper searches in the IRAS data base with different strategies \\citep{auma:85, walk:88, mann:98}. For a review see \\citet{back:93}.  Recently \\citet{plet:99} have discussed these earlier results and concluded that a significant excess at $60\\mu$m is found in $13\\pm 10\\%$ of all main sequence stars with spectral type A, F, G and K.  Unfortunately all these studies based on IRAS data only were affected by severe selection effects and did not answer important questions such as: Will a star loose its disk when it grows older? On what time-scale?  Does the presence of planets depend on the stellar main-sequence mass?  Do multiple stars have disks more, or less frequently? Do stars that formed in clusters have disks less often?  With such questions unanswered we clearly do not understand the systematics of the formation of solar systems.  Here we present results of a continuation with ISO \\citep{kess:96} of the succesful search of IRAS. Our aim has been to obtain a better defined sample of stars. The major step forward in this paper is not in the detection of more remnant disks, but in reliable information about the presence \\textit{or absence} of a disk.  Earlier reports on results from our program have been given in \\citet{habi:96}, \\citet{domi:98}, \\citet{jour:99} and \\citet{habi:99}. ",
        "conclusions": " \\begin{itemize} \\item Fourteen stars have a flux in excess of the expected photospheric flux. We conclude that each of these has a circumstellar disk that is a remnant from its pre-main-sequence time. Two more stars may have a disk, but there is a significant chance that the emission is due to a background galaxy. \\item The overall incidence of disks is 14/84 or 17\\%. A-stars have a higher incidence than the other stars. \\item We prove that the detectability of a given disk is the same for A, F and G-stars with the same photospheric flux at 60\\um ; K stars have a lower probability of detection. \\item The disks that we detect have a value of $\\tau$, between $2\\times 10^{-4}$ and $4\\times 10^{-6}$. The upper limit is real: main-sequence stars do not carry stronger disks; the exception is $\\beta$ Pic with $\\tau=8\\times 10^{-3}$. The lower limit is caused by selection effects: fainter disks are below our detection threshold. \\item Six out of the ten A-type stars younger than 400 Myr have a disk; the disk is absent around all five older A-type stars: the disks disappear around this age. \\item Disks around F-, G- and K-stars probably disappear on a similar time scale. The disappearance of disks is \\textit{not} a continuous process; 400 Myr is the age at which most disks disappear. \\item In the history of our solar system the abrupt ending of the initial ``heavy bombardment'' of the Moon has the same time scale. \\item We suggest that the disks that we detect are actually sites where a ``heavy bombardment'' takes place now. The time scale on which the disks disappear is actually the time scale of the disappearance of the bombarding planetesimals. \\item The mass that we detect through its infrared emission is only a minute fraction (about $10^{-5}$) of the mass present. We see the gravel but not the very big stones that produce it. \\item Some very young stars lack a disk; some very old stars still have a disk: the existence of both groups needs to be explained. \\item Stars in multiple systems retain their remnant disks as often as isolated stars. \\end{itemize}"
    },
    "0011/astro-ph0011301_arXiv.txt": {
        "abstract": "We calculate the energy distribution of the relativistic particles injected in the ICM during a phase of reacceleration of the relativistic particles in the cluster volume. We apply our results to the case of the Coma cluster in which recent merging activity and the presence of the radio halo may suggest that reacceleration processes are efficient. We find that the electron population injected in the central part of the cluster by the head--tail radio galaxy NGC 4869 may account for a large fraction, if not all, of the detected EUV excess via inverse Compton scattering of the cosmic microwave background (CMB) photons.  If radio haloes are powered by reacceleration mechanisms active in the cluster volume, moderate non--thermal EUV excesses (of order of $\\sim 1-5 \\cdot 10^{42}$erg s$^{-1}$) should be a common feature of clusters containing powerful head--tail radio galaxies and/or AGNs.   \\medskip \\par\\noindent {\\it PACS}: 95.30.Cq; 95.30.Gv; 98.54.Gr; 98.65.Cw; 98.65.Hb ",
        "introduction": "Observations of a few galaxy clusters with the {\\it Extreme Ultraviolet Explorer} have revealed extreme ultraviolet emission (EUV) in excess of that expected by extrapolating downward the thermal X--ray spectrum (Lieu et al. 1996a,b; Bowyer, Lampton and Lieu 1996; Mittaz, Lieu and Lockman 1998). The EUV excess has been claimed by these authors as a new feature of clusters of galaxies with important physical implications.  The suggestion that spurious EUV excesses may derive by underestimating the line--of--sight Galactic absorption (Arabadjis \\& Bregman 1999) has been excluded by further EUV investigations (Bowyer, Berg\\\"ofer and Korpela 1999). These authors have also pointed out that the measured EUVs are strongly influenced by the variation of the telescope sensitivity over the field of view. By making use of the appropriate corrections, EUV extended excess was confirmed in the case of Coma (Bowyer et al. 1999) and Virgo (Bergh\\\"{o}fer, Bowyer and Korpela 2000), whereas no evidence of extended EUV excess was found in A1795 and A2199 (Bowyer et al. 1999). The presence of EUV excess in the latter two clusters is, however, still debated (Bonamente, Lieu and Mittaz 2000; Lieu, Bonamente and Mittaz 2000).  The EUV excess may be interpreted as due to a relatively cool ($\\sim 10^6$ K) emitting gas (e.g. Lieu et al.1996a), but, at these temperatures the gas cooling is particularly efficient, so that a non--thermal interpretation is usually favoured with respect to the thermal scenario (Hwang 1997; Ensslin \\& Biermann 1998; Bowyer \\& Bergh\\\"ofer 1998; Sarazin \\& Lieu; Lieu et al. 1999). Indeed, synchrotron emission in the radio band and inverse Compton (IC) emission in the EUV and hard X--rays are expected from clusters of galaxies in the framework of continuous injection of relativistic electrons (Sarazin 1999; V\\\"{o}lk \\& Atoyan 1999) and of a {\\it two phase} model invoking reacceleration of previously injected relic particles (Brunetti et al. 2000).  Due to the large amount of data the Coma cluster represents the best case to compare model predictions with observations. Since the EUV profile is narrower than the radio one, it cannot be accounted for by IC emission of the low energy halo electrons with the CMB photons (Bowyer \\& Bergh\\\"{o}fer 1998). Thus, alternative scenarios have been explored.  Ensslin et al. (1999) have calculated that anisotropic IC scattering of the optical photons (from the central galaxies in the cluster) by an assumed ad hoc very dense population of mildly relativistic electrons (with a cut--off in the energy distribution at a few MeV) can, in principle, account for the EUV properties.   Atoyan \\& V\\\"{o}lk (2000) have argued that the EUV can be accounted for by IC scattering of the CMB photons from a relic electron population with a radiative cut--off at $\\gamma \\geq 500$, injected in the ICM at the epoch of starburst activity and recently slightly re--energized.  Sarazin (1999; 2000) has pointed out that clusters of galaxies can contain a large number of relic relativistic electrons (with $\\gamma < 300$) injected in the past which can account for the EUV excess via IC scattering of CMB photons.  In a recent paper Brunetti et al.(2000) have shown that the radio spectral steepening observed in the case of Coma C (Giovannini et al. 1993; Deiss et al. 1997) may imply the presence of diffuse reacceleration in the cluster volume and the spectrum of a relic population would be stretched toward higher energies thus reducing the emitted EUV flux.  Following the original proposal of Brunetti et al.(1999), in this paper we investigate the possibility that a substantial fraction, if not all, of the EUV excess emission from the Coma cluster can be accounted for by the IC scattering of the CMB photons due to relativistic electrons recently injected in the central part of the cluster by AGN activity and reaccelerated by moderate turbulence in the ICM.  In Sections 2.1 and 2.2 we illustrate the proposed scenario and in Section 2.3 we discuss the possibility that NGC 4869 may supply the required number of relativistic particles; a general discussion is given in Section 3. In Appendix A we calculate the time evolution of the energy distribution of relativistic electrons continuously injected and reaccelerated in the ICM, while relevant formulae used to calculate the synchrotron and IC emitted spectra are given in the Appendix B.  $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ is assumed throughout. ",
        "conclusions": "It has been previously shown (Brunetti et al 2000) that the radio (synchrotron) and hard X-rays (IC on the CMB photons) properties of Coma C may be explained by a two phase model in which diffuse relic electrons have been efficiently reaccelerated throughout the cluster volume for a period of $\\sim$1 Gyr by the turbulence and associated Fermi type mechanisms excited in the ICM by a recent merger (main electron population, {\\it MeP}).  By adopting this model we have shown here that a steady injection in the cluster core of an additional fresh population of relativistic electrons (Fig.6) may explain the EUV excess, recently determined by Bowyer et al.(1999), as IC scattering of the CMB photons. Detailed computations of the time evolution of the injected electron spectrum show that internal consistency within the complete model requires a relatively steep injection spectral index ($\\delta \\sim 2.6$), however typical of the radio galaxies, and an injection starting time $\\leq 0.3$ Gyr, in which case there is no significant contribution to the synchrotron emission in the 327-1400 MHz band and to the total X-ray flux. Nevertheless, it should be pointed out that an important contribution to the emission at low radio frequencies ($\\leq 100$ MHz; Fig.4) may be present, so that improved measurements at long radio wavelengths may provide a test of the proposed interpretation of the EUV excess and, of course, constrain the model parameters.  The luminosity of the EUV excess is $\\sim 1.5 \\cdot 10^{42}$ erg/sec. An IC origin with the spectral shape of our model needs a total number of injected electrons $\\sim 4\\cdot 10^{63}$. We have explored the possibility that a large fraction of these electrons may have been injected by the powerful head tail radio galaxy NGC 4869 orbiting in the cluster core. By assuming that NGC 4869 has been constantly active in the past 0.3 Gyr (an age comparable to that of a number of giant and head tail radio galaxies) we find, based on minimum energy conditions in the tail, that about 40\\% of the required number of electrons may have been released to the ICM; a full complement is achieved with a moderate departure from the equipartition resulting in a magnetic field intensity about a factor $\\sim 3$ weaker. The total energy of the electrons deposited in the ICM is $\\sim 4 \\cdot 10^{59}$ erg which is not unreasonable for a radio galaxy.  We stress that, although the internal pressure is considerably increased with respect to previous estimates based on classical equipartition formulae, the radio tail is still well confined by the hot ICM: the pressure gap is reduced by almost an order of magnitude, but the tail is still underpressurized by a factor 2--3. The introduction of a filling factor and of a relativistic proton component may now easily bring to a pressure equilibrium.  In Fig.7 the observed EUV radial profile is compared with that predicted by our model based on the simple assumption of a spherical and uniform distribution of the relativistic particles centered on the cluster centre. It is seen that the model provides a very good fit of the data at an angular distance $> 3$ arcmin from the centre which includes $\\sim 90-95\\%$ of the EUV excess. We note that in the innermost 3 arcmin the radial profile samples a region containing the cD galaxy NGC 4874 so that an additional contribution to the EUV flux may be expected. Furthermore, it should be stressed that at smaller radii the azimuthal distribution of the EUV excess is very sensitive to small scale brightness/temperature fluctuations of the hot ICM. If the relativistic electrons are mainly supplied by NGC 4869 then the EUV brightness distribution should be related to some extent with the orbital path followed by the tailed radio source. At present the tail's end is positioned at an angular distance of $\\sim 6$ arcmin from the cluster centre, so that a shift between the X-ray and EUV maxima may be expected. However, most of the injected relativistic plasma could have been stirred around the cluster core in the available time scale by the turbulence generated by the random motion of the galaxies; this may have considerably smoothed out the distribution of the relativistic particles on the cluster core size and thus that of the emitted EUV photons. Detailed mapping of the EUV excess would obviously be of great importance.  We point out that our model, although specifically developed for the Coma cluster, may have a wider application predicting the existence of sizeable fluxes of excess EUV photons from galaxy clusters where large numbers of relativistic electrons are injected in the cluster cores by powerful AGN activity, or other energetic events such as a minor merger, during a reacceleration phase. However, due to the requirement of a contemporaneous presence of efficient Fermi type turbulence reacceleration and of an efficient recent injection of relativistic electrons in the cluster core (i.e. the presence of powerful AGNs), we claim that in the clusters with radio haloes a non--thermal EUV excess due to the proposed scenario would be more rare than the hard X--ray excess.  Finally, if a relic population of relativistic electrons is commonly associated with clusters of galaxies (Sarazin 1999), the related EUV excess from IC scattering of CMB photons is expected to be more common in clusters without radio halo than in those with radio halo where diffuse reacceleration processes may stretch toward higher energies the electron spectrum.   \\begin{figure} \\special{psfile=brunetti_f6.ps hoffset=-30 voffset=-420 vscale=60 hscale=60} \\vspace{12 cm} \\caption{ The energy distribution integrated over the cluster volume of the main electron population (dashed line) and of the additional electron population (solid lines) are shown in arbitrary units. The typical energies of the electron emitting via IC in the hard X--ray band (HEX) and in the EUV band are also showed.} \\end{figure}  \\begin{figure} \\special{psfile=brunetti_f7.ps hoffset=-30 voffset=-420 vscale=60 hscale=60} \\vspace{12 cm} \\caption{The comparison between the model calculations and the observed EUV profile is reported. The model EUV profile (solid line) is obtained by taking into account the instrumental background plus the expected EUV from the X--ray plasma (dashed line). The data for the instrumental background and EUV from X--ray plasma and observed EUV excess (bold data points) are taken from Bowyer et al. (1999).} \\end{figure}   \\begin{ack} We are grateful to S. Bowyer and T. Ensslin for useful discussions and suggestions. This work was partly supported by the Italian Ministry for University and Research (MURST) under grant Cofin98-02-32. \\end{ack}"
    },
    "0011/astro-ph0011184_arXiv.txt": {
        "abstract": "We examine the characteristics of nucleosynthesis in 'hypernovae', i.e., supernovae with very large explosion energies ($ \\gsim 10^{52} $ ergs).  We carry out detailed nucleosynthesis calculations for these energetic explosions and compare the yields with those of ordinary core-collapse supernovae.  We find that both complete and incomplete Si-burning takes place over more extended, lower density regions, so that the alpha-rich freezeout is enhanced and produces more Ti in comparison with ordinary supernova nucleosynthesis.  In addition, oxygen and carbon burning takes place in more extended, lower density regions than in ordinary supernovae.  Therefore, the fuel elements O, C, Al are less abundant while a larger amount of Si, S, Ar, and Ca (\"Si\") are synthesized by oxygen burning; this leads to larger ratios of \"Si\"/O in the ejecta.  Enhancement of the mass ratio between complete and incomplete Si-burning regions in the ejecta may explain the abundance ratios among iron-peak elements in metal-poor stars. Also the enhanced \"Si\"/O ratio may explain the abundance ratios observed in star burst galaxies.  We also discuss other implications of enhanced [Ti/Fe] and [Fe/O] for Galactic chemical evolution and the abundances of low mass black hole binaries. ",
        "introduction": "\\label{sec:intro}  Massive stars in the range of 8 to $\\sim$ 150$M_\\odot$ undergo core-collapse at the end of their evolution and become Type II and Ib/c supernovae unless the entire star collapses into a black hole with no mass ejection (e.g., Arnett 1996).  These Type II and Ib/c supernovae (as well as Type Ia supernovae) release large explosion energies and eject explosive nucleosynthesis materials, thus having strong dynamical, thermal, and chemical influences on the evolution of interstellar matter and galaxies.  Therefore, the explosion energies of core-collapse supernovae are fundamentally important quantities, and an estimate of $E \\sim 1\\times 10^{51}$ ergs has often been used in calculating nucleosynthesis (e.g., Woosley \\& Weaver 1995; Thielemann et al. 1996) and the impact on the interstellar medium. (In the present paper, we use the explosion energy $E$ for the final kinetic energy of explosion.) A good example of this estimate is SN1987A in the Large Magellanic Cloud, whose energy is estimated to be $E = (1.0$ - 1.5) $\\times$ $10^{51}$ ergs from its early light curve (Shigeyama et al. 1987, 1988; Woosley et al. 1988; Arnett et al. 1989; Nakamura et al. 1998; Blinnikov et al. 2000).   SN1998bw called into question the applicability of the above energy estimate for all core-collapse supernovae.  This supernova was discovered in the error box of the gamma-ray burst GRB980425 (Soffitta et al. 1998; Lidman et al. 1998; Galama et al. 1998).  Although the association of supernovae with gamma-ray bursts is still the subject of controversy (e.g., Wang \\& Wheeler 1998; Norris et al. 1999), SN1998bw is a quite unusual supernova in its radio and optical properties (Galama et al. 1998; Iwamoto et al. 1998; Kulkarni et al. 1998).  It is one of the most luminous supernovae at radio wavelengths and its peak luminosity was attained exceptionally early in comparison with other radio supernovae (Kulkarni et al. 1998). From the optical spectra SN1998bw is classified as a Type Ic supernova (SN Ic), but it shows unusually broad spectral features.  The modeling of the early light curve and spectra leads us to conclude that it was an explosion of a $\\sim$ 14 $M_\\odot$ C + O star with the explosion energy $E$ = 3 - 6 $ \\times$ $10^{52}$ ergs, which is about thirty times larger than that of a canonical supernova (Iwamoto et al. 1998; Woosley, Eastman \\& Schmidt 1999; Branch 2000; Nakamura et al. 2000, 2001).  The amount of $^{56}$Ni ejected from SN1998bw is found to be $M(^{56}$Ni) $\\simeq$ 0.4 - 0.7 $M_\\odot$ (Nakamura et al. 2000, 2001; Sollerman et al. 2000), which is about 10 times larger than the 0.07$M_\\odot$ produced in SN1987A, a typical value for core-collapse supernovae.  We have used the term 'hypernova' to describe such an extremely energetic supernova with $E \\gsim 10^{52}$ ergs (Nomoto et al. 2000). In terms of explosion energies, pair-instability supernovae of 150 - 200 $M_\\odot$ stars also explode with $E \\gsim 10^{52}$ ergs, thus being called hypernovae (e.g.,Ober et al. 1983; Bond et al. 1984; Stringfellow \\& Woosley 1988). In the present paper, we concentrate on nucleosynthesis in core-collapse hypernovae.  Recently, other hypernovae candidates have been recognized.  SN1997ef and SN1998ey are also classified as SNe Ic, and show very broad spectral features similar to SN1998bw (Garnavich et al. 1997, 1998).  The spectra and the light curve of SN1997ef have been well simulated by the explosion of a 10$M_\\odot$ C+O star with $E = 1.0 \\pm 0.2$ $\\times$ $10^{52}$ ergs and $M(^{56}$Ni) $\\sim$ 0.15 $M_\\odot$ (Iwamoto et al. 2000; Mazzali, Iwamoto, \\& Nomoto 2000).  SN1997cy is classified as a SN IIn and unusually bright (Germany et al. 2000; Turatto et al. 2000). Its light curve has been simulated by a circumstellar interaction model which requires $E \\sim$ 5 $\\times$ $10^{52}$ ergs (Turatto et al. 2000).  The spectral similarity of SN1999E to SN1997cy (Cappellaro et al. 1999) would suggest that SN1999E is also a hypernova.  Note that all of these estimates of $E$ assume a spherically symmetric event.  In this paper, we explore nucleosynthesis in such energetic core-collapse induced supernovae, the systematic study of which has not yet been done. In \\S \\ref{sec:98bw}, we briefly describe the argument that the progenitor mass $M$ and explosion energy $E$ are very large from the observation of SN1998bw.  In \\S \\ref{sec:model} and \\S \\ref{sec:depend}, the characteristics of nucleosynthesis in hypernovae are investigated with detailed nucleosynthesis calculations and compared with nucleosynthesis in canonical supernovae.  We also discuss possible effects on the Galactic chemical evolution and on the abundances in metal-poor stars in \\S \\ref{sec:diss}. A summary is given in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We investigated explosive nucleosynthesis in hypernovae, that is, hyper-energetic supernovae. Detailed nucleosynthesis calculations were performed and compared to those of ordinary core-collapse supernovae.  We also studied implications to Galactic chemical evolution and the abundances in metal-poor stars.  We demonstrated, using the SN1998bw data, that progenitor of hypernovae must be massive based on the information from light curve and nucleosynthesis calculations of $^{56}$Ni, without using spectral information. This argument is useful for cases where reliable spectral information is not available.  In hypernovae, both complete and incomplete Si-burning takes place in more extended, and hence, lower density regions, so that the alpha-rich freezeout is enhanced in comparison with normal supernova nucleosynthesis.  Thus $^{44}$Ca, $^{48}$Ti, and $^{64}$Zn are produced more abundantly than in canonical supernovae.  Oxygen and carbon burning also takes place in more extended regions for the larger explosion energy, hence the yield of these fuel elements are smaller than canonical supernovae.  We found that hypernova nucleosynthetic yields are compatible with the observed Cr, Mn, and Fe abundances of metal-poor halo stars, provided that the explosion energy is $\\sim$ 10$^{52}$ ergs. Our results imply that heavier progenitors allow larger explosion energy to match the abundance pattern of metal-poor stars, although further investigations are necessary for firm conclusion. Our results suggest that some metal-poor halo stars may bear the abundance pattern of a single hypernova.  We also suggest that the problem of the deficit of Ti in the Galactic chemical evolution with the supernova yields currently available could be settled by the contribution of hypernovae, in which $\\alpha$-rich freezeout is enhanced and large amounts of $\\alpha$-elements are ejected.  The abundance pattern of the starburst galaxy M82, characterized by abundant Si and S relative to O and Fe, may be attributed to hypernova explosions if the IMF is relatively flat, and thus the contribution of massive stars to the galactic chemical evolution is large.  The black hole binary X-ray Nova Sco (GRO J1655-40) may also be affected by a hypernova explosion, which can eject large amount of Ti, S, and Si and leave a massive black hole as a compact remnant.  In this paper, we performed calculations assuming spherical symmetry and thus did not take account of non-spherical effects.  For example, if a jet is associated with hypernovae, the explosive shock may be pointed and stronger in this direction.  In this case, the high explosion energies along the jet should have the properties of nucleosynthesis in hypernovae discussed above and orthogonal to the jet may be more like ordinary supernovae. In such conditions, a hypernova could make ejecta which have a range of abundance ratios (as seen in Figure 6) in different directions. If those ejecta in various directions interact with interstellar matter and form mixed materials of various mass ratios, stars which could form from such interactions would have a range of metallicity and abundance ratios among iron-peak elements.  For example, the ejecta in the jet direction has larger amount of complete Si-burning products (i.e., larger [Co/Fe]) and are mixed with larger amount of hydrogen (i.e., producing smaller [Fe/H]) because of higher velocities than in other directions.  Then such explosions might be responsible for the abundance trends in the metal-poor stars.  Multi-dimensional simulations are needed to investigate these possibilities.  \\bigskip  We would like to thank Drs. B. Nordstr\\\"om, F. Spite, T. Tsuru, G. Israelian, G. Brown, Ph. Podsiadlowski, P. Mazzali, and B. Schmidt for stimulating discussion on the observational indication of hypernova nucleosynthesis. This work has been supported in part by the Grant-in-Aid for Scientific Research (12640233, 12740122) and COE research (07CE2002) of the Japanese Ministry of Education, Science, Culture, and Sports, Swiss National Science Foundation grant 2000-53798.98, U.S. National Aeronautics and Space Administration grant NAG5-8405 and the Joint Institute for Heavy Ion Research, which has as member institutions the University of Tennessee, Vanderbilt University, and the Oak Ridge National Laboratory.  ORNL is managed by UT-Battelle, LLC, for the U.S. Department of Energy under contract DE-AC05-00OR22725."
    },
    "0011/astro-ph0011467_arXiv.txt": {
        "abstract": "We construct a starburst model for the hidden starbursts in Arp 220 based on the new Starburst99 models of Leitherer et al. Comparing these stellar population synthesis models with observations, we show that the  hidden power source must be due to star formation (as opposed to an AGN) at the 50\\% level or more in order to avoid an ionizing photon excess problem, and this starburst must be young ($< 7 \\times 10^7$ yr). We derive a current star formation rate of $270 M_{\\odot}$ yr$^{-1}$, and an extinction $A_V > 30$ mag for our line of sight to this hidden starburst. ",
        "introduction": "Ultraluminous infrared galaxies (ULIGs) are the most luminous galaxies in the Universe and their bolometric luminosities are comparable to those of quasars ($L_{\\rm bol} > 10^{12}L_{\\odot}$). The physical origin of the large luminosity of ULIGs is unknown since the very centers of these objects, which is where the bolometric luminosity comes from (Sanders \\& Mirabel 1996) is very heavily obscured by dust and direct observation is not possible. Either an enshrouded starburst or an active galactic nucleus (AGN), or both, could generate the observed luminoisites.  Arp 220 is the nearest and best-studied ULIG\\footnote{We use a distance of 74 Mpc, assuming $H_0=75 \\; {\\rm km \\; s^{-1} \\; Mpc^{-1}}$ and $\\Omega_0=1$ with the recession velocity to the Galactic Standard of Rest, $V_{\\rm GSR}=5,531 {\\rm km \\; s^{-1}}$ (de Vaucouleurs et al. 1991).} Recently Genzel et al.~(1998) performed a mid-infrared spectroscopy program for a sample of ULIGs and demonstrated that Arp 220 has an appreciable PAH feature and so is powered by a starburst, assuming that the very core where $L_{\\rm bol }$ originates from is not optically thick at 15 $\\mu$m. Based on detailed radiative transfer modelling, Rowan-Robinson (2000 and references therein) argues that galaxies like Arp 220 must be star-formation powered since they emit most of their luminosity longward of 50$\\mu$m. Additional indirect evidence that Arp 220 is powered by star formation includes \\vskip 1pt \\noindent (i) the lack of an X-ray detection (Iwasawa 1999), even at hard wavelengths where Compton scattering is less important; \\vskip 1pt \\noindent (ii) VLBI radio observations, which show many radio supernovae but no strong point radio source which could correspond to the AGN (Smith et al.~1998); \\vskip 1pt \\noindent (iii) a LINER\\footnote{LINER = Low Ionization Nuclear Emission-line Region (Heckman 1980).}-like optical emission-line spectrum which can result from a starburst-driven superwind (Veilleux et al. 1995; Taniguchi et al. 1999; Lutz, Veilleux, \\& Genzel 1999). This is not, however, unambiguous evidence that there is no AGN present. In any case, these measurements do not probe the very central regions, which is where most of the power is emitted.  Rather they probe regions with optical depths $\\tau < 1$ at optical wavelengths.  \\vskip 1pt \\noindent On the other hand OH megamaser emission certainly indicates that the power source is compact (Lonsdale et al.~1994, 1998). This would be consistent with an AGN power source, but again this measurement is not a unique discriminant. The nature of the power source in cold ULIGs like Arp 220 has taken on a particular significance in recent times with the detection of a population of these objects at high redshifts which contribute a very large fraction of the extragalactic background light at far-infrared and submillimeter wavelengths (Barger, Cowie \\& Sanders 1999).  We now attempt to provide additional constraints on the nature of the power source in Arp 220 by investigating the so-called ionizing photon deficiency problem. Prestwich et al.~(1994) performed near-infrared (NIR) spectroscopy for a sample of starburst galaxies including Arp 220 and found that the ratio of the Br$\\gamma$ luminosity to the infrared luminosity of Arp 220 is lower by a factor of ten than what is observed in lower-luminosity (but still very luminous) starburst galaxies, suggesting a deficit of ionizing photons. This could in part be due to absorption of the ionizing photons by dust, to an AGN contributing much of $L_{\\rm bol}$ but few ionizing photons, or to the detailed shape of the stellar initial mass function (IMF). These physical process all operate on the global properties of Arp 220 in different ways and in this paper we use data from the literature together with starburst models and an SED fitting method to give constraints on the nature of the hidden power source. ",
        "conclusions": "In all the cases that we consider, the hidden power source due to star formation at the 50\\% level or more in order to avoid an ionizing photon excess problem, and this starburst must be young ($< 7 \\times 10^7$ yr). We derive a current star formation rate of $267 M_{\\odot}$ yr$^{-1}$, and an extinction $A_V > 30$ mag to this hidden starburst.  \\vspace{1pc} YS thanks the Japan Society for Promotion of Science (JSPS) Research Fellowships for Young Scientist. NT thanks the PPARC for financial support. YT would like to thank R. -P. Kudritzki, Bob McLaren and Dave Sanders at Institute for Astronomy, University of Hawaii for their warm hospitality. This work was financially supported in part by Grant-in-Aids for the Scientific Research (Nos.\\ 07044054, 10044052, and 10304013) of the Japanese Ministry of Education, Science, Sports and Culture."
    },
    "0011/gr-qc0011085_arXiv.txt": {
        "abstract": "We consider the problem of searching for continuous gravitational wave (CW) sources orbiting  a companion object. This issue is of particular interest because, the Low Mass X-ray Binaries (LMXB's), and among them Sco X-1 -- the brightest X-ray source in the sky -- might be marginally detectable with $\\approx 2$ years coherent observation time by the Earth-based laser interferometers expected to come on line by 2002, and clearly observable by the second generation of detectors. Moreover, several radio pulsars, which could be deemed to be CW sources, are found to orbit a companion star or planet, and the LIGO/VIRGO/GEO600 network plans to continuously monitor such systems.  We estimate the computational costs for a search launched over the additional five parameters describing generic elliptical orbits (up to $e\\simlt 0.8$) using match filtering techniques. These techniques provide the optimal signal-to-noise ratio and also a very clear and transparent theoretical framework. Since matched filtering will be implemented in the final and the most computationally expensive stage of the hierarchical strategies, the theoretical framework provided here can be used to determine the computational costs. In order to disentangle the computational burden involved in the orbital motion of the CW source, from the other source parameters (position in the sky and spin-down), and reduce the complexity of the analysis, we assume that the source is monochromatic (there is no intrinsic change in its frequency) and its location in the sky is exactly known. The orbital elements, on the other hand, are either assumed to be completely unknown or only partly known.  We provide ready-to-use analytical expressions for the number of templates required to carry out the searches in the astrophysically relevant regions of the parameter space, and how the computational cost scales with the ranges of the parameters. We also determine the critical accuracy to which a particular parameter must be known, so that no search is needed for it; we provide rigorous statements, based on the geometrical formulation of data analysis, concerning the size of the parameter space so that a particular neutron star is a one-filter target. This result is formulated in a completely general form, independent of the particular kind of source, and can be applied to any class of signals whose waveform can be accurately predicted.  We apply our theoretical analysis to Sco X-1 and the 44 neutron stars with binary companions which are listed in the most updated version of the radio pulsar catalogue. For up to $\\approx 3$ hours of coherent integration time, Sco X-1 will need at most a few templates; for a week's integration time the number of templates rapidly rises to $\\simeq 5\\times 10^6$. This is due to the rather poor measurements available today of the projected semi-major axis and the orbital phase of the neutron star. If, however, the same search is to be carried out with only a few filters, then more refined measurements of the orbital parameters are called for - an improvement of about three orders of magnitude in the accuracy is required. Further, we show that the five NS's (radio pulsars) for which the upper-limits on the signal strength are highest, require no more than a few templates each and can be targeted very cheaply in terms of CPU time.  Blind searches of the parameter space of orbital elements are, in general, completely un-affordable for present or near future dedicated computational resources, when the coherent integration time is of the order of the orbital period or longer. For wide binary systems, when the observation covers only a fraction of one orbit, the computational burden reduces enormously, and becomes affordable for a significant region of the parameter space. ",
        "introduction": "\\label{sec:intro}  The construction of several large-scale interferometric gravitational wave detectors, with optimal sensitivity in the frequency window $\\sim$ 10 Hz -- 1 kHz is close to completion. For a period extending upto a year, starting from the end of year 2000, engineering data runs will be carried out in order to test and debug the detector components.  Finally, between 2002 and 2004 the interferometers will carry out the first set of \"science\" data-runs with the {\\it realistic goal of  directly observing gravitational waves} (GW's). This initial phase will be followed by substantial upgrades on most of the instruments aimed at reaching  better sensitivity and  larger observational bands. The projects include the Laser Interferometer Gravitational-wave Observatory (LIGO) in the USA consisting of two facilities, one at Hanford (WA) and the other at Livingston (LA), hosting two 4-km and one 2-km interferometers~\\cite{ligo}; VIRGO, a French/Italian project located at Cascina, near Pisa, Italy, which consists of a 3-km detector (and currently running about one-to-two years behind the time-frame presented above) ~\\cite{virgo}; GEO600, a German/British 0.6-km interferometer located at Ruthe, near Hannover, Germany; in Japan, the TAMA project is {\\it currently running} a medium scale interferometric detector of arm length 300 meters, and is planning to extend the baseline to 3 km and carry out other substantial improvements on the instrument within the next few years~\\cite{tama}; finally, the ACIGA consortium will build, if funding is approved, a 500 meter interferometer (AIGO500) near Perth, Australia~\\cite{aigo}. In the meantime, a number of existing resonant bar detectors are steadily increasing their sensitivity in narrow bandwidths ($\\sim 1$ Hz) covering the kHz spectral window~\\cite{Cerdonio} and serious efforts are being made to fly, possibly by 2010, a space-borne laser interferometer~\\cite{lisa} (LISA: the Laser Interferometer Space Antenna). LISA will open the low frequency window $10^{-5}$ Hz -- $10^{-2}$ Hz, currently accessible only via the Doppler tracking of interplanetary space-craft~\\cite{BVI99}, and is projected to have sensitivity several orders of magnitude better.  Building and running such powerful machines represents an enormous enterprise, which has started over 30 years ago. However, analyzing the large amount of data -- several Mbytes per second -- and digging out with high confidence astrophysical signals from the noise which severely corrupts the data, presents its own challenges. In fact, over the past few years, interest has been growing in the area of GW data analysis, which is now regarded as one of the key aspects for the successful detection of GWs.  Several types of GW sources have been envisaged which could be directly observed by Earth-based detectors (see~\\cite{Thorne87,Thorne95,Fl_GR15,Schutz99} and references therein for recent reviews): (i) Burst sources -- such as binary systems of neutron stars (NS) and/or black holes (BH) in their in-spiral phase, BH/BH and/or BH/NS mergers, and supernovae explosions -- whose signals last for a time much shorter, typically between a few milli-seconds and a few minutes, than the planned observational time; (ii) stochastic backgrounds of radiation, either of primordial or astrophysical origin, and (iii) continuous wave (CW) sources -- e.g. rapidly rotating neutron stars -- where a weak deterministic signal is continuously present in the data stream.  In this paper we deal with data analysis issues regarding CW sources. More specifically, we investigate the computational load involved in filtering the data stream to search for {\\it monochromatic radiation emitted by a neutron star orbiting a companion object}.  \\subsection{The data analysis challenge} \\label{subs:dan}  CW emitters pose one of the most computationally intensive problems in GW data analysis. The weakness of the expected signal requires very long observation times, of the order of a year (or possibly more) for accumulating enough signal-to-noise ratio (SNR) so that detection is possible. During the observation time, a monochromatic signal in the source reference frame, is Doppler modulated by the motion of the detector carried by the spinning Earth orbiting the Sun. The emitted energy is spread over $\\simeq 2\\times 10^6\\,(T/10^7\\,{\\rm sec})^2\\, (f/1\\,{\\rm kHz})$ frequency bins of width $\\Delta f = 1/T$, where $T$ is the time of observation (the formula holds for $T$ upto six months, after that the number of bins increases linearly with $T$). In order to recover the whole power in one frequency bin one has to \"correct\" the recorded data stream for each possible source position in the sky. The problem is made worse, if the the intrinsic frequency of the source changes, say due to spindown. Then the power is spread over $3\\times 10^6\\,(\\tau/10^3\\,{\\rm yrs})^{-1}\\,(T/10^7\\,{\\rm sec})^2\\,(f/1\\,{\\rm kHz})$ bins, where $\\tau = f/\\dot{f}$ is (twice) the spin-down age of the NS. Indeed, one then needs to correct, in addition, for this effect searching through one or more spin-down parameters. It is clear that searches for CW's are limited by the available computational resources~\\cite{Schutz89,BCCS98}. The optimal technique of tracking coherently the signal phase -- matched filtering -- is not affordable for integration times longer than $\\sim$ 1 to 10 days (the integration time depends on the range of parameters over which the search is being carried out). So far, the only viable strategy for a \"blind\" search over a wide range of parameters is based on a hierarchical structure~\\cite{BC00,SP99}. In this type of search, coherent and incoherent search stages are alternated in order to identify candidate signals -- and the relevant parameter range -- \"cheaply\" by a suboptimal algorithm. Then, the candidate signals are followed-up with a coherent and computationally expensive search over the entire observation time -- this step constitutes the bottle neck for processing power -- to recover the optimal SNR.  However, due to the large computational burden, the algorithms investigated so far have been restricted to {\\it isolated} NS's. The problem of searching for a CW emitter orbiting a companion object is simply considered computationally intractable, as it would require upto five more search parameters which would compound the already enormous computational cost.  To gain insight into the additional Doppler modulation introduced on the signal phase by the source motion around a companion, let us consider a NS, of mass $m_{\\rm NS} = 1.4\\Ms$, with an orbiting companion object of mass $m_2$ and period $P$. During the observation time $T$, the NS velocity changes by $\\Delta v = (2\\pi/P)^{4/3}\\,(G m_{\\rm NS})^{1/3}\\,q/(1+q)^{2/3}\\,T$, where $q \\equiv m_2/m_{\\rm NS}$. The frequency is Doppler shifted by $\\Delta_D\\,f = \\pm f \\Delta v/c$ which has to be compared to the frequency resolution  $\\Delta f$. (Here and after, $G$ denotes the Newton's gravitational constant and $c$ denotes the speed of light or actually GW). The maximum observational time $T_{{\\rm max}}$ after which one needs to correct for the CW's source orbital motion is therefore, \\be T_{{\\rm max}} \\simeq 131.6 \\, \\left[\\frac{(1+q)^{1/3}}{\\sqrt{q}}\\right]\\, \\left(\\frac{P}{1\\,{\\rm day}}\\right)^{2/3}\\, \\left(\\frac{m_{NS}}{1.4\\,\\Ms}\\right)^{-1/6}\\, \\left(\\frac{f}{1\\,{\\rm kHz}}\\right)^{-1/2}\\,{\\rm sec}\\,. \\label{Tmax} \\ee Of course, this time scale can vary by orders of magnitude, depending on the NS companion and the orbital period. For the relevant astrophysical situations we have: \\be T_{{\\rm max}} \\simeq \\left\\{ \\begin{array}{ll} 1.2\\times 10^5\\,\\left(\\frac{P}{10^7\\,{\\rm sec}}\\right)^{2/3}\\, \\left(\\frac{m_{2}}{10^{-3}\\,\\Ms}\\right)^{-1/2}\\, \\left(\\frac{f}{1\\,{\\rm kHz}}\\right)^{-1/2}\\,{\\rm sec} & (m_2 \\ll m_{\\rm NS}) \\nonumber\\\\ 1.0\\times 10^2\\,\\left(\\frac{P}{12\\,{\\rm hours}}\\right)^{2/3}\\, \\left(\\frac{f}{1\\,{\\rm kHz}}\\right)^{-1/2}\\,{\\rm sec} & (m_2 \\simeq m_{\\rm NS})\\nonumber\\\\ 4.4\\times 10^2\\,\\left(\\frac{P}{10\\,{\\rm days}}\\right)^{2/3}\\, \\left(\\frac{m_{2}}{10\\,\\Ms}\\right)^{-1/6}\\, \\left(\\frac{f}{1\\,{\\rm kHz}}\\right)^{-1/2}\\,{\\rm sec} & (m_2 \\gg m_{\\rm NS})\\\\ \\end{array} \\right. \\,; \\label{Tmax1} \\ee For comparison, we note that the maximum integration time before which the motion of the detector from Earth's rotation about its axis, produces a Doppler shift in frequency greater than the frequency resolution bin is $T_{\\rm max}^{\\rm (det)} \\simeq 50\\, (f/1\\,{\\rm kHz})^{-1/2}$ min. It is then evident that in the case of a CW source orbiting a companion object, the frequency shift is typically more severe and the computational challenge that one faces when searching for the additional 5-dimensional parameter vector describing the binary orbit is considerable.  \\subsection{Expected sources and motivation} \\label{subs:sm}  The {\\it characteristic amplitude} of the continuous signal emitted by a triaxial NS, rotating about a principal axis with period $P_r$ and frequency $f_r = 1/P_r$ reads~\\cite{Thorne87}: \\be h_{\\rm c} \\simeq 7.7 \\times 10^{-26}\\, \\left(\\frac{\\epsilon_e}{10^{-6}}\\right)\\, \\left(\\frac{D}{10\\,{\\rm kpc}}\\right)^{-1}\\, \\left(\\frac{f_r}{500\\,{\\rm Hz}}\\right)^{2} , \\label{hchar} \\ee where $\\epsilon_e$ is the equatorial ellipticity of the star, and we have taken the standard value $10^{45}$ gr cm$^2$ for the NS moment of inertia. In Eq.~(\\ref{hchar}) we assume that the GW energy is released at exactly twice the rotation frequency, so that $f = 2 f_r$.  The characteristic amplitude has to be compared with the noise fluctuations, \\be h_{\\rm n} = \\sqrt{\\frac{S_n(f)}{T}} = 10^{-25}\\, \\left(\\frac{S_n(f)}{10^{-43}\\,{\\rm Hz}^{-1}}\\right)^{1/2}\\, \\left(\\frac{T}{10^7\\,{\\rm sec}}\\right)^{-1/2}\\, \\label{hn} \\ee in a resolution bin, where $S_n(f)$ is the noise spectral density of the instrument. A signal is detected if $h_{\\rm c} \\simgt k h_{\\rm n}$, where the multiplicative factor $k$, of order of a few, depends on the number of filters used in the analysis and on the (subjective) confidence level of the measurement. The strength of the signal depends crucially on the $\\epsilon_e$, which is a measure of the non-axis-symmetry of a star. Although large uncertainties exist in the maximum value that $\\epsilon_e$ can achieve, theoretical studies of the braking strain of NS crusts suggest that $\\epsilon_e \\simlt 10^{-5}$. One can observationally constrain $\\epsilon_e$ if the NS rotational period $P_r$ and its first time derivative $\\dot{P_r}$ are known: \\be \\epsilon_e \\le 5.1 \\times 10^{-7}\\, \\left(\\frac{P_r}{2\\,{\\rm msec}}\\right)^{3/2}\\, \\left(\\frac{\\dot{P_r}}{10^{-15}}\\right)^{1/2}\\, . \\label{ell} \\ee This yields an upper-limit on the amplitude of the emitted GW's, by assuming that all energy is lost through gravitational wave emission. By combining Eqs.~(\\ref{hchar}) and~(\\ref{ell}), we can immediately place the upper-limit \\be h_{\\rm c} \\simlt 3.9\\times 10^{-26}\\, \\left(\\frac{f_r}{500\\,{\\rm Hz}}\\right)^{1/2}\\, \\left(\\frac{\\dot{P_r}}{10^{-15}}\\right)^{1/2}\\, \\left(\\frac{D}{10\\,{\\rm kpc}}\\right)^{-1}\\, . \\label{hchar1} \\ee The comparison of $h_{\\rm c}$ for some known NS's with $h_{\\rm n}$ of the laser interferometers is shown in Fig.~\\ref{fig:h}.  \\begin{figure} \\begin{center} \\epsfig{file=binaryns_fig1.eps,width=4.1in} \\caption{\\label{fig:h} Upper limits on the GW amplitude from known radio pulsars and Sco X-1. The plot shows the upper-limit on $h_{\\rm c}$, Eq.~(\\protect{\\ref{hchar1}}), for the neutron stars included in the latest version of the published radio pulsar catalogue~{\\protect\\cite{cat}}, for which the spin-down and the distance are known (isolated pulsars: squares; binary pulsars: bold bullets), and the estimated value of $h_{\\rm c}$ for Sco X-1 (triangle), Eq.~(\\protect{\\ref{hchar2}}). We also show the noise amplitude $h_{\\rm n}$ for an observation time $T = 10^7$ sec, for the first generation detectors (GEO600: solid line; LIGO I: dotted line; VIRGO: dashed line).} \\end{center} \\end{figure}%  In Table~\\ref{tab:pulsars} we also show some of the parameters of the binary radio pulsars which are known so far. It is interesting to notice that a bunch of systems are within the reach of the first generation of instruments for coherent observation times from 15 months to 4 years, and well within the enhanced configuration sensitivity. It is also worth noticing that they represent a fair sample of orbital companions, as they include NS/NS binaries, NS/white-dwarf binaries, NS/main-sequence stars and NS/planet. However, the actual signal strain could be orders of magnitude below the upper-limit~(\\ref{hchar1}).  A far more promising scenario is the accretion of hot material onto the NS surface. Here the induced quadrupole moment is directly related to the accretion rate, which can be copious. The gravitational energy reservoir, moreover, can be continuously replenished, if persistent accretion occurs. The key idea behind this scenario is that gravitational wave radiation can balance the torque due to accretion, and was proposed over 20 years ago~\\cite{PP78,Wagoner84}. However, it has attracted considerable new interest in the past two years and has been fully revitalized  by the launch of the Rossi X-ray Timing Explorer, designed for precision timing of accreting NS's. % \\vbox{ \\begin{table} \\begin{center} \\caption{\\label{tab:pulsars} Neutron stars detected as radio pulsars in binary systems. The table shows the main parameters of the 44 known neutron stars detected as radio pulsars orbiting a companion object listed in the Taylor et al. catalogue~{\\protect\\cite{cat}}. We list the gravitational wave frequency $f$ -- assumed to be exactly twice the radio frequency -- its first time derivative $\\dot{f}$, the estimated distance $D$ from the solar system and the spin-down age (we do not give all significant digits). The orbital parameters -- the orbital period $P$, the projection of the semi-major axis along the line of sight $a_{\\rm p}$ and the eccentricity $e$ -- are tabulated with their respective errors (in parenthesis). Notice that out of the 44 known systems, seven sources radiate at a frequency $< 10$ Hz and therefore are not observable by the (so far) planned ground-based experiments. } \\begin{tabular}{ll|cccc|ccc} PSR J & PSR B & $f$      & $\\dot{f}$              & $D$ & $\\log({\\rm age}/{\\rm yr})$ & $P$  & $a_{\\rm p}$ & $e$ \\\\ &       & ${\\rm Hz}$ & $10^{-15}\\,{\\rm sec}^{-2}$ & kpc &                       & days & sec              &     \\\\ \\hline 0024-7204E   & 0021-72E &  565.6   &    --    &  4.5  &  --     &    2.256844 (3)       &     1.978 (3)      &   0.000 (3) \\cr 0024-7204I   & 0021-72I &  573.9   &    --    &  4.5  &  --     &    0.226 (3)          &     0.039 (2)      &   0.00 (5) \\cr 0024-7204J   & 0021-72J &  952.1   &    --    &  4.5  &  --     &    0.120665 (2)       &     0.0402 (2)     &   0.00 (5) \\cr 0034-0534    &          &  1065.4  &     3.8  &  1.0  &    9.6  &    1.58928180 (3)     &     1.437768 (5)   &   0.0000 (1) \\cr 0045-7319    &          &  2.2     &    10.5  & 57.0  &    6.5  &   51.169226 (3)       &   174.2540 (8)     &   0.807995 (5) \\\\ 0218+4232    &          &  860.9   &    27.8  &  5.8  &    8.7  &    2.02885 (1)        &     1.9844 (1)     &   0.00000 (2) \\\\ 0437-4715    &          &  347.4   &     3.4  &  0.1  &    9.2  &    5.741042329 (12)   &     3.3666787 (14) &   0.0000187 (10) \\cr 0613-0200    &          &  653.2   &     2.3  &  2.2  &    9.6  &    1.19851251 (2)     &     1.09145 (2)    &   0.000000 (22) \\\\ 0700+6418    & 0655+64  &  10.2    &     0.04 &  0.5  &    9.7  &    1.028669703 (1)    &     4.125612 (5)   &   0.0000075 (11) \\cr 0751+1807    &          &  574.9   &     1.3  &  2.0  &    9.8  &    0.2631442679 (4)   &     0.396615 (2)   &   0.0000 (1) \\\\ 0823+0159    & 0820+02  &  2.3     &     0.3  &  1.4  &    8.1  & 1232.47 (12)          &   162.1466 (7)     &   0.011868 (7) \\cr 1012+5307    &          &  380.5   &     1.1  &  0.5  &    9.8  &    0.604672713 (5)    &     0.581816 (6)   &   0.00000 (2) \\\\ 1022+10      &          &  121.6   &    --    &  0.6  &  --     &    7.80513014 (8)     &    16.765412 (6)   &   0.0000977 (6) \\cr 1045-4509    &          &  267.6   &     0.7  &  3.2  &    9.8  &    4.0835291 (1)      &     3.015107 (9)   &   0.000019 (6) \\cr 1300+1240    & 1257+12  &  321.6   &     5.9  &  0.6  &    8.9  &   66.536 (1)          &     0.0013106 (6)  &   0.0182 (9) \\cr 1302-6350    & 1259-63  &  41.9    &  1994.5  &  4.6  &    5.5  & 1236.72359 (5)        &  1296.580 (2)      &   0.869931 (1)   \\cr 1312+1810    & 1310+18  &  60.3    &    --    & 18.9  &  --     &  255.8 (6)            &    84.2 (7)        &   0.002 (3) \\cr 1455-3330    &          &  250.4   &    --    &  0.7  &  --     &   76.17458 (2)        &    32.36222 (6)    &   0.000167 (2) \\cr 1518+4904    &          &  48.9    &    --    &  0.7  &  --     &    8.6340047 (8)      &    20.043982 (8)   &   0.2494848 (7)    \\cr 1537+1155    & 1534+12  &  52.8    &     3.4  &  0.7  &    8.4  &    0.4207372998 (3)   &     3.729468 (9)   &   0.2736779 (6)   \\cr 1623-2631    & 1620-26  &  180.6   &    12.9  &  1.8  &    8.3  &  191.442819 (2)       &    64.809476 (13)  &   0.0253147 (5) \\cr 1640+2224    &          &  632.2   &     0.6  &  1.2  &   10.2  &  175.460668 (2)       &    55.329722 (2)   &   0.00079742 (5) \\cr 1641+3627B   & 1639+36B &  566.9   &    --    &  7.7  &  --     &    1.259113 (3)       &     1.389 (2)      &   0.005 (5) \\cr 1643-1224    &          &  432.7   &     3.1  &  4.9  &    9.3  &  147.01743 (6)        &    25.07260 (4)    &   0.000506 (2) \\cr 1713+0747    &          &  437.6   &     0.8  &  1.1  &    9.9  &   67.82512988 (2)     &    32.342413 (1)   &   0.00007492 (2) \\cr 1721-1936    & 1718-19  &  2.0     &     3.2  & 11.6  &    7.0  &    0.2582735 (2)      &     0.3526 (8)     &   0.000 (5) \\\\ 1748-2446A   & 1744-24A &  173.0   &    -0.3  &  7.1  &  --     &    0.0756461170 (4)   &     0.11963 (2)    &   0.0000 (12) \\\\ 1803-2712    & 1800-27  &    6.0   &     0.3  &  3.6  &    8.5  &  406.781 (2)          &    58.9397 (18)    &   0.000507 (6) \\cr 1804-0735    & 1802-07  &  86.6    &     1.8  &  3.2  &    8.9  &    2.6167634 (5)      &     3.92047 (4)    &   0.211999 (15)   \\cr 1804-2718    &          &  214.1   &    --    &  1.2  &  --     &   11.15 (1)           &     7.20 (4)       &   0.00 (1) \\\\ 1823-1115    & 1820-11  &  7.1     &    35.2  &  6.3  &    6.5  &  357.7622 (3)         &   200.672 (3)      &   0.79462 (1)    \\cr 1834-0010    & 1831-00  &  3.8     &     0.1  &  2.6  &    8.8  &    1.811103 (2)       &     0.7231 (8)     &   0.000 (4) \\\\ 1857+0943    & 1855+09  &  373.0   &     1.2  &  0.9  &    9.7  &   12.3271711905 (6)   &     9.2307802 (4)  &   0.00002168 (5) \\cr 1910+0004    &          &  552.7   &    --    &  4.1  &  --     &    0.140996 (1)       &     0.038 (2)      &   0.00 (1) \\\\ 1915+1606    & 1913+16  &  33.9    &     4.9  &  7.1  &    8.0  &    0.322997462736 (7) &     2.3417592 (19) &   0.6171308 (4)    \\cr 1955+2908    & 1953+29  &  326.1   &     1.6  &  5.4  &    9.5  &  117.349097 (3)       &    31.412686 (5)   &   0.0003304 (3) \\cr 1959+2048    & 1957+20  &  1244.2  &    13.0  &  1.5  &    9.2  &    0.3819666389 (13)  &     0.892268 (8)   &   0.00000 (4) \\\\ 2019+2425    &          &  508.3   &     0.9  &  0.9  &    9.9  &   76.5116347 (1)      &    38.767625 (1)   &   0.00011111 (6)\\cr 2033+17      &          &  336.2   &    --    &  1.4  &  --     &   56.2 (1)            &    20.07 (8)       &   0.00 (5) \\\\ 2130+1210C   & 2127+11C &  65.5    &    10.7  & 10.0  &    8.0  &    0.335282052 (6)    &     2.520 (3)      &   0.68141 (2)    \\cr 2145-0750    &          &  124.6   &    --    &  0.5  &  --     &    6.83890256 (8)     &    10.16411 (1)    &   0.000021 (2) \\\\ 2229+2643    &          &  671.6   &     0.4  &  1.4  &   10.4  &   93.015895 (2)       &    18.912519 (5)   &   0.0002550 (4) \\cr 2305+4707    & 2303+46  &  1.9     &     1.0  &  4.4  &    7.5  &   12.33954454 (17)    &    32.6878 (3)     &   0.658369 (9)    \\cr 2317+1439    &          &  580.5   &     0.4  &  1.9  &   10.4  &    2.459331464 (2)    &     2.3139483 (9)  &   0.0000005 (7) \\cr \\end{tabular} \\end{center} \\end{table}% } The observational evidence that Low Mass X-ray Binaries (LMXB's) - binary systems where a compact object accretes material from a low mass companion -- in our Galaxy are clustered around a rotation frequency $\\approx 300$ Hz, led Bildsten~\\cite{Bildsten98} to propose a mechanism to explain this behaviour. The fundamental idea is that continuous emission of GW's radiates away the angular momentum that is transfered to the NS by the infalling material. The fact that the rate of angular momentum loss through GW's scales as $f^5$, provides a very natural justification of the clustering of rotation frequency of several sources. The physical process responsible for producing a net quadrupole moment is the change of composition in the NS crust, which in turn is produced by the temperature gradient caused by the in-falling hot material. Recently, Ushomirsky et al.~\\cite{UCB00} have posed this initial idea on more solid theoretical grounds. If such mechanism does operate, LMXB's are extremely interesting candidate sources for Earth-based detectors. Several systems would be detectable by LIGO operating in the \"enhanced\" configuration (LIGO II), if the detector sensitivity is tuned, through narrow-banding, around the emission frequency. In particular, Sco X-1, the most luminous X-ray source in the sky, possibly is marginally detectable by \"initial\" LIGO and GEO600 (the latter in narrow-band configuration), where an integration time of approximately 2 years would be required. The characteristic amplitude for this class of sources is: \\be h_{\\rm c} \\simeq 4\\times 10^{-27}\\,\\left(\\frac{R_{NS}}{10^6\\,{\\rm cm}}\\right)^{3/4}\\, \\left(\\frac{m_{NS}}{1.4\\,\\Ms}\\right)^{-1/4}\\, \\left(\\frac{F_X}{10^{-8}\\,{\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm sec}^{-1}}\\right)^{-1/2} , \\label{hchar2} \\ee where $R_{NS}$ is the typical radius of a neutron star and $F_X$ is the X-ray flux.  Data analysis strategies to search for NS's in binary systems might therefore be vital for ground-based detectors. They would supplement the already existing strategies for the several sources mentioned above.   \\subsection{Organization of the paper}   The paper is organized as follows: In Section \\ref{sec:model} we set up the parametrized model of the signal based on the two body problem. We then identify the search parameters of the signal, which are essentially the binary orbital elements, for the general case of elliptic orbits, as well as for circular orbits. In Section \\ref{sec:datan} we discuss the geometrical approach to the data analysis; in Section~\\ref{subs:nfilt} we review the method of determining the number of templates which involves the computation of the metric on the signal manifold, and we also state useful properties of the metric which are then used in the later sections. In Section~\\ref{subs:tar} we present a rigorous criterion, naturally suggested by the geometrical picture, for estimating the number of filters required in targeted searches and to determine whether any particular source is a one-filter target.  Sections~\\ref{sec:cost} and~\\ref{sec:obs} contain the key results of the paper. In Section~\\ref{sec:cost} we obtain analytic approximate formulae for the number density of filters per unit coordinate volume in the parameter space in the two important regimes of observation time: (i) observations lasting for several orbital periods, and (ii) observations lasting a fraction of an orbit. These expressions provide for us the scaling: that is, how the computational cost scales with different parameters. They can also be readily applied to estimate the computational cost of detecting GW from some observed systems. Further, we also compare these results with the fully numerical computation of the metric determinant (which shows excellent agreement over a wide range of the parameters), and determine the range of applicability of these approximate expressions as a function of the relevant parameters (in particular the time of observation as compared to the orbital period). In the Section~\\ref{sec:obs} we apply the above analysis to known NS's in binary systems, specifically radio-pulsars and LMXB's. Given the current (or near future) computational resources, we then determine the systems that could be searched by matched filtering methods. Section~\\ref{sec:concl} contains our final remarks and future directions of work. ",
        "conclusions": "\\label{sec:concl}  We have estimated the extra-computational costs required to search for CW sources in binary systems. This work was motivated by the fact that LMXB's, a key class of sources, possibly, for the initial generation of detectors and certainly for the enhanced detectors, are NS's orbiting a binary companion. Also several NS's detected in the radio band are in binaries and will be continously monitored by the Earth-based network of GW detectors. In order to disentangle the computational costs arising from correcting for the Doppler phase modulation produced by the source motion around the orbital companion, we have assumed that the signal is exactly monochromatic and the location of the source in the sky is perfectly known. Our results provide an estimate of the computational burden associated with these extra dimensions of the parameter space which has been neglected so far (for rather obvious reasons). We have deduced closed form analytical expressions for the number of templates required to carry out the analysis. We have obtained the {\\it scaling} of the computational costs as a function of the orbital elements and the observational parameters (the source emission frequency and the observation time); they are ready-to-use tools to evaluate the computational resources needed and the trade off for search strategies. We have applied our analysis to known radio pulsars and Sco X-1, estimating the computational costs involved in targeting these systems. The analysis also addresses the need of further astronomical observations of known NS's, with the aim of reducing the errors in the source parameters, so that the search is made inexpensive. In fact, a byproduct of our analysis is the rigorous formulation, in the geometrical framework of data analysis, of the conditions, {\\it i.e.} the size of the parameter space, such that a NS can be considered a one-filter target; our result is completely general, and can be applied to any class of sources (not necessarily CW sources) whose waveforms can be modelled reliably.  Our analysis is clearly limited in several respects. Possibly, the main accomplishment is a quantitative understanding of the key issues that require intensive attack, and identifying the bottlenecks that prevent us from carrying out blind searches of NS in binary systems. There are three clear directions of future work that we believe should be pursued vigorously, particularly in the light of the current data analysis effort in setting up GW search codes:  \\begin{enumerate}  \\item The software implementation of matched filter based searches of a NS in a binary system, for given values of the source parameters such as location in the sky, spin-down(s), orbital elements. This task is simple in principle,  as the fundamental building blocks are already available: search codes for GW's from isolated NS's and publicly available software (such as TEMPO) in the radio pulsar community to correct for the Doppler effect of the signal phase for a source in a binary system. However, the implementation of this might not prove to be so trivial.  \\item The estimation of the computational costs keeping into account {\\it both} the orbital elements {\\it and} the the source position and the spin-down parameters. This is the composite problem where the full expression for the total costs needs to be obtained. It will take into account possible correlations between the search parameters investigated here, and those that describe isolated neutron stars.  \\item The investigation of suitable {\\it hierarchical strategies}, that can reduce the computational costs as compared to fully coherent searches. The current vigorous effort in the study, and software implementation, of hierarchical algorithms for isolated sources, the so-called stack-and-slide and Hough-pattern-tracking searches provide already a template for investigations for a larger parameter space. Moreover, radio and X-ray astronomers are currently investigating algorithms suitable to carry out analogous searches in the electro-magnetic spectrum~\\cite{JK91,Ransom00}.  \\end{enumerate}  One limitation of our approach is that we have completely neglected relativistic effects in the model of a binary orbit. In particular, the advance of the periapse is very significant in relativistic systems (the largest periapse advance reported in~\\cite{cat} is $\\approx 4$ deg/yr). This, and other complications, with respect to the simple Keplerian model that we have adopted here, need to be taken into account in future studies, and addressed carefully.  We would like to conclude by stressing that the data analysis problem that we have addressed here is common to electro-magnetic observations of pulsating stars in binary orbits, where typical examples are radio and X-rays binaries. We can therefore hope that the current efforts in other related research fields can help in providing a breakthrough in this area."
    },
    "0011/astro-ph0011431_arXiv.txt": {
        "abstract": "{ The XMM-Newton first-light image revealed X-ray point sources which show heavily absorbed power-law spectra. The spectral indices and the probable identification of a radio counterpart for the brightest source suggest AGN shining through the interstellar gas of the Large Magellanic Cloud (LMC). The column densities derived from the X-ray spectra in combination with H\\,{\\sc i} measurements will allow to draw conclusions on H\\,{\\sc i} to H$_2$ ratios in the LMC and compare these with values found for the galactic plane. ",
        "introduction": "The study of active galactic nuclei and quasi stellar objects (both named AGN hereafter) behind nearby galaxies allows to probe the interstellar matter in the foreground galaxy. Absorption lines in optical spectra provide information about the absorbing material along the line of sight through the galaxy. However, very few AGN are known behind nearby galaxies as surveys for these objects avoided the crowded fields. Dedicated searches for AGN behind nearby galaxies were performed by Tinney et al. (\\cite{tcz97}) and  Crampton et al. (\\cite{cgc97}). Conducting optical identification programs of ROSAT X-ray sources, they found 15 background objects in the direction of the Small Magellanic Cloud and 11 towards the LMC involving also clusters of galaxies and individual distant galaxies. For all these background objects red-shifts were determined covering the range of 0.02 -- 2.3. From more general optical identification programs of X-ray sources in the Magellanic Clouds by Cowley et al. (\\cite{cch84,csm97}) and Schmidtke et al. (\\cite{scc99}) further AGN were found and X-ray classification work by Haberl \\& Pietsch (\\cite{hp99}) and Haberl et al. (\\cite{hfp00}) suggests candidates for background objects based on their X-ray (and radio) properties. However, as consequence of the strong attenuation of X-rays in the ROSAT 0.1 -- 2.4 keV band by interstellar matter which hampers the X-ray detection of AGN in the dense cores (where column densities of the order of \\ohcm{22} are expected which suppress the flux at 1 keV by a factor $\\sim$10 and at 0.4 keV by more than 10$^4$), the background objects were mainly found in the outer regions of the Magellanic Clouds.  The high sensitivity and good spectral resolution of the EPIC instruments on board of XMM-Newton (Jansen et al. \\cite{j01}) over a wide energy band from 0.1 keV up to 15 keV provides a unique tool to detect hard (intrinsically and/or highly absorbed) X-ray sources and allows for the first time to separate very accurately the intrinsic X-ray source spectrum (determined from higher energies) and the photo-electric absorption which attenuates the spectrum at lower energies. The EPIC instruments are therefore ideally suited to investigate the absorbing matter in the dense cores of nearby galaxies by observing AGN behind the galaxies. While radio observations in the 21 cm line infer the column density of H\\,{\\sc i} only (e.g. Arabadijs \\& Bregman \\cite{ab99}), X-ray absorption measurements provide a linear measure of the total mass along the line of sight.  In this letter we report on a sample of highly absorbed point sources detected in the XMM-Newton first-light image obtained by the EPIC-PN camera (Str\\\"uder et al. \\cite{s01}) south-west of the 30 Doradus region in the LMC. Their X-ray spectra strongly suggest AGN as origin of the X-rays. ",
        "conclusions": "We investigated nine point sources with hard X-ray spectra detected in the XMM-Newton first-light images. Their X-ray spectra can be represented by highly absorbed power-laws and the derived photon indices between 1.7 and 2.0 are typical for steep AGN spectra (e.g. Vaughan et al. \\cite{vrw99}). In contrast accreting high mass X-ray binaries show harder X-ray spectra with typical indices between 0 and 1 (e.g. Yokogawa et al. \\cite{yit00}) not compatible with the indices found for our point sources. The radio spectral index of 0.73 found for \\object{J\\xmc} also supports the AGN interpretation.  AGN spectra sometimes show intrinsic absorption which can easily reach the amount seen in the spectra reported here. So the question remains which fraction of the column density can be attributed to the LMC. The fact that the three close point sources near the center of the field of view all show similar high absorption of the order of \\ohcm{22} suggests that most of the \\nh\\ is intrinsic to the LMC. The derived column density of 0.35\\hcm{22} of sources only $\\sim$2\\arcmin\\ away from the central sources shows that the absorption can change by a factor of more than 5 over distances of $\\sim$30 pc in the LMC. It is remarkable in this respect that the north-west part of the first-light image lacks of point sources. This suggests that even higher column densities exceeding \\ohcm{22} are required to suppress the detection of AGN in that region.  X-ray absorption measurements are directly related to the total hydrogen column density, comprising H\\,{\\sc i}, H$_2$ and H\\,{\\sc ii} column densities. Because the absorbing elements at X-ray energies are He and metals their relative abundance must be taken into account to derive X-ray hydrogen column densities. This was done in our spectral fits assuming abundances of 0.5 relative to solar (Russell \\& Dopita \\cite{rd92}). Kim et al. (\\cite{kss00}) have produced a H\\,{\\sc i} column density map of the LMC with a resolution of 1\\arcmin\\ which shows variations from 0 to 5\\hcm{21} with maximum values around the 30\\,Doradus region. The higher absorption seen in the XMM sources suggests either large amounts of H\\,{\\sc ii} and/or H$_2$ gas or higher metal abundances along their lines of sight through the LMC. Higher abundances of O, Ne, Mg and Si in the interstellar medium are suggested by the spectrum of the diffuse emission near the 30 Dor region (Dennerl et al. \\cite{dha01})."
    },
    "0011/astro-ph0011577_arXiv.txt": {
        "abstract": " ",
        "introduction": "As will be apparent to anyone reading this book, the practitioners of N-body simulations have an enormous variety of preoccupations.  Some are essentially pure mathematicians, who view the field as an exciting application for abstruse theory.  Others enjoy formulating and tackling mathematically-neat problems, with little concern over whether the particular restrictions that they impose are also respected by nature.  Still others are closer to computer scientists, inspired by the challenge of developing ever more sophisticated algorithms to tackle the N-body problem, but showing less interest in the ultimate application of their codes to solving astrophysical problems.  This contribution is presented from yet another biased perspective: that of the observational galactic dynamicist. Observational astronomers tend to use N-body simulations in a rather cavalier manner, both as a tool for interpreting existing astronomical data, and as a powerful technique by which new observations can be motivated.  The aim of this article is to illustrate this profitable interplay between simulations and observations in the study of galaxy dynamics, as well as highlighting a few of its shortcomings.  To this end, the text of this chapter is laid out as follows. Section~\\ref{kindynsec} provides an introduction to the sorts of data that can be obtained in order to study the dynamical properties of galaxies, and goes on to discuss the intrinsic stellar dynamics that one is trying to model with these observations, and the role that N-body simulations can play in this modeling process. Section~\\ref{nbodyhistsec} gives a brief overview of the historical development of N-body simulations as a tool for studying galaxy dynamics.  Section~\\ref{ellipticalsec} provides some examples of the interplay between observations and N-body simulations in the study of elliptical galaxies, while Section~\\ref{disksec} provides further examples from studies of disk galaxies, concentrating on barred systems.  These sections are in no way intended to be encyclopedic in scope; rather, by selecting a few examples and examining them in some detail, the text seeks to give some flavour for the range of what is possible in this rich field.  Finally, Section~\\ref{futuresec} contains some speculations as to what may lie in the future for this productive relationship between observations and N-body modeling. ",
        "conclusions": ""
    },
    "0011/astro-ph0011094_arXiv.txt": {
        "abstract": "The basic assumptions behind Population Synthesis and Spectral Evolution models are reviewed. The numerical problems encountered by the standard population synthesis technique when applied to models with truncated star formation rates are described. The Isochrone Synthesis algorithm is introduced as a means to circumvent these problems. A summary of results from the application of this algorithm to model galaxy spectra by Bruzual and Charlot (1993, 2000) follows. I present a comparison of these population synthesis model predictions with observed spectra and color magnitude diagrams for stellar systems of various ages and metallicities. It is argued that models built using different ingredients differ in the resulting values of some basic quantities (e.g. $M/L_V$), without need to invoking violations of physical principles. The range of allowed colors in the observer frame is explored for several galaxy redshifts. \\footnote{To appear in {\\it Proceedings of the XI Canary Islands Winter School of Astrophysics on Galaxies at High Redshift}, eds. I. P\\'erez-Fournon, M. Balcells and F. S\\'anchez} ",
        "introduction": "The number distribution of the stellar populations present in a galaxy is a function of time. Thus, the number of stars of a given spectral type, luminosity class, and metallicity content changes as the galaxy ages. In early-type galaxies (E/S0) most of the stars were formed during, or very early after the initial collapse of the galaxy and the stellar population ages as times goes by. The chemical abundance in these systems must have reached the value measured in the stars very quickly during the formation process since most E/S0 galaxies show little evidence of recent major events of star formation. In late-type galaxies the stellar population also ages, but there is a significant number of new stars being formed. Depending on the star formation rate, $\\Psi(t)$, the mean age of the stars in a galaxy may even decrease as the galaxy gets older. In general, in late-type systems the metal content of the stars and the interstellar medium is an increasing function of time. $\\Psi(t)$ can also increase above its typical value due to interactions between two or more galaxies or with the environment.  As a consequence of the aging of the stellar population, or its renewal in the case of galaxies with high recent $\\Psi(t)$, the rest-frame spectral distribution of the light emitted by a galaxy is a function of its proper time. Observational properties such as photometric magnitude and colors, line strength indices, metal content of gas and young stars, depend on the epoch at which we observe a galaxy on its reference frame. This {\\it intrinsic} evolution should not be mistaken with the {\\it apparent} evolution produced by the cosmological redshift $z$. For distant galaxies both effects may be equally significant.  In order to search for spectral evolution in galaxy samples we must {\\it (a)} quantify the amount of evolution expected between cosmological epoch $t_1$ and $t_2$, and {\\it (b)} design observational tests that will reveal this amount of evolution, if present. Evolutionary population synthesis models predict the amount of evolution expected under different scenarios and allow us to judge the feasibility of measuring it. In the ideal case (Arag\\'on-Salamanca et al. 1993; Stanford et al. 1995, 1998; Bender et al. 1996) spectral evolution is measured simply by comparing spectra of galaxies obtained in such a way that the same rest frame wavelength region is sampled in all galaxies, irrespective of $z$. In this case there is no need to apply uncertain $K$-corrections to transform all the spectra to a common wavelength scale. Alternatively, if this approach cannot be applied, e.g. when studying large samples of faint galaxies, one can use models which allow for different degrees of evolution (including none) to derive indirectly the amount of evolution consistent with the data (Pozzetti et al. 1996; Metcalfe et al. 1996). Clearly, the first approach is to be preferred whenever possible. In all cases we must rule out possible deviations from the natural or passive evolution of the stellar population in some of the galaxies under study induced, for example, by interactions with other galaxies or cluster environment, etc.  The large amount of astrophysical data that has become available in the last few years has made possible to build several complete sets of stellar population synthesis models. The predictions of these models have been used to study many types of stellar systems, from local normal galaxies to the most distant galaxies discovered so far (approaching $z$ of 4 to 5), from globular clusters in our galaxy to proto-globular clusters forming in different environments in distant, interacting galaxies. In this paper I present an overview of results from population synthesis models directly applicable to the interpretation of galaxy spectra. ",
        "conclusions": "Present population synthesis models show reasonable agreement with the observed spectrum of stellar populations of various ages and metal content. Differences in results from different codes can be understood in terms of the different ingredients used to build the models and do not necessarily represent violations of physical principles by some of these models. However, inspection of Fig. 20 shows that two different sets of evolutionary tracks for stars of the same metallicity produce models that at early ages differ in brightness and color from 0.5 to 1 mag, depending on the specific bands. The differences decrease at present ages in the rest frame, but are large in the observer frame at $z > 2$. Thus any attempt to date distant galaxies, for instance, based on fitting observed colors to these lines will produce ages that depend critically on the set of models which is used. Note that from $z$ of 3 to 3.5 $(V-K)$ in the two models differs by more than 1 mag. This difference is produced by the corresponding difference between the models seen in the rest frame at 10 Myr. From Figs. 15 and 16 these differences can be understood in terms of the different contribution of the same stellar groups to the total $V$ and $K$ flux in the two models.  Even though at the present age models built with different IMFs show reasonably similar colors and brightness, the early evolution of these models is quite different at early ages (Fig. 21), resulting in larger color differences in the observer frame at $z > 2$. Thus, the more we know about the IMF, the better the model predictions can be constrained. The small color differences seen in the rest frame when different stellar libraries of the same metallicity are used, are magnified in the observer frame (Fig. 22). When the $k$ correction brings opposing flux differences into each filter, the difference in the resulting color is enhanced. Fig. 23 shows the danger of interpreting data for one stellar system with models of the wrong metallicity. The color differences between these models, especially in the observer frame, are so large as to make any conclusion thus derived very uncertain.  It is common practice to use solar metallicity models when no information is available about the chemical abundance of a given stellar system. Galaxies evolving according the $Z(t)$ laws of Fig. 24 show color differences with respect to the $Z=Z_\\odot$ model which are not larger than the differences introduced by the other sources of uncertainties discussed so far. Hence, the solar metallicity approximation may be justified in some instances. The color differences between the chemically inhomogeneous composite population and the purely solar case (Fig. 25), are much smaller than the ones shown in Fig. 23 for chemically homogeneous SSPs.  Figs. 26 to 39 indicate that some colors, especially $(V-R)$, when measured in the observer frame are less sensitive to model predictions than other colors. From Figs. 26 to 31, metallicity $Z$ and the SFR are the most dominant factors determining the range of allowed colors.  I expect that through these simple examples the reader can get a feeling of the kind of uncertainties introduced by the many ingredients entering the stellar population synthesis problem, and that he or she will be motivated to try his or her own error estimates when using these models."
    },
    "0011/astro-ph0011027_arXiv.txt": {
        "abstract": "We report the serendipitous discovery of a flare star observed with the ROSAT X-ray observatory.  From optical spectra, which show strong and variable emission lines of the hydrogen Balmer series and neutral helium, we classify this object as a M3.0Ve star, and estimate a distance of 52 pc from published photometry.  Due to the star's close proximity (13.6\\arcmin) to the calibration source and RS CVn binary AR Lacertae, long term X-ray coverage is available in the ROSAT archive ($\\sim$50 hours spanning 6.5 years). Two large flare events occurred early in the mission (6-7/1990), and the end of a third flare was detected in 6/1996.  One flare, observed with the Position Sensitive Proportional Counter (PSPC), had a peak luminosity L$_{X}$=1.1$\\times10^{30}$ erg s$^{-1}$, an e-folding rise time of 2.2 hours and a decay time of 7 hours. This decay time is one of the longest detected on a dMe star, providing evidence for the possibility of additional heating during the decay phase.  A large HRI flare (peak L$_{X}$=2.9$\\times10^{30}$ erg s$^{-1}$) is also studied.  The ``background\" X-ray emission is also variable - evidence for low-level flaring or microflaring.  We find that $\\geq$59\\% of the HRI counts and $\\geq$68\\% of the PSPC counts are due to flares.  At least 41\\% of the HRI exposure time and 47\\% of the PSPC are affected by detectable flare enhancement. ",
        "introduction": "X-ray emission from late-type M dwarfs has been studied extensively to investigate the structure and emission mechanisms of stellar coronae. Coronal heating to X-ray emitting temperatures is attributed to either impulsive flares or quiescent energy release in magnetic structures. The corona of these stars are thought to be similar to the sun but often with luminosities orders of magnitude higher.  The high magnetic activity of these flare stars is also seen in their optical spectra. Continuum enhancement and strong emission lines of the H Balmer series, Ca II and neutral He are often evident (Montes et al., 1999).  An EXOSAT study (Pallavicini et al. 1990) showed flares have a wide range of energies and timescales.  Most outbursts can be described as either implusive (decay time $<$ 1 hr) or long decay flares (decay time $>$ 1 hr) and have thermal X-ray spectra with temperatures similar to solar X-ray flares.  The impulsive stellar flares have timescales similar to the compact solar flares. The long duration flares have greater total energy and are more similar to two-ribbon flare events.  From ROSAT observations, coronal emission from dMe stars has been shown to have two distinct spectral components, a low temperature component attributed to quiescent active regions and a variable, high temperature component due to compact flaring regions (Giampapa et al. 1996).  Schmitt (1994) has shown conclusively the existence of long duration flares on M stars using the ROSAT all-sky survey.  In some flaring stars, the long decay can be attributed to continual heating of the flaring region (Schmitt \\& Favata, 1999; Ottmann \\& Schmitt, 1996). New flare models have been developed (Reale et al., 1997) in which the additional heating determines the characteristics of the decay. Using this model, an analysis of long duration flares on AD Leo (Favata, Micela \\& Reale, 2000) and EV Lac (Favata et al., 2000), have shown the emitting regions to be compact with length scales less than the stellar radius and similar in size to solar flares,  thereby providing evidence that long duration flares are produced in high pressure structures.   During an analysis of observations with the ROSAT High Resolution Imager (HRI) of the RS CVn binary AR Lac, we noticed an X-ray source within the field of view to be highly variable and undetected in many fields.  The first X-ray detection of this source was with the $\\textit{Einstein}$ observatory  (hence the catalogue name, 2E 2206.6+4517; Harris et al., 1993).  Coaddition of six IPC observations spanning 26.6 ksec yield a 4$\\sigma$ detection with 115 net counts. The source is undetected in a 1.5 ksec HRI observation.  No flaring activity is evident during these observations.  We analyzed thirty nine observations from the ROSAT public data archive (16 PSPC; 23 HRI) which included this source position within the field of view.  The total observing time was $\\approx$50 hours. The observations span two flares with almost complete light curves, the tail end of a third flare, and show variability in the low-level X-ray emission.  The flare detected with the PSPC has a long decay time, possibly providing evidence for significant continual heating during the flare decay.  Thus, a further in-depth study of flares and quiescent emission from 2E 2206.6+4517 could provide a useful test of current models. ",
        "conclusions": "In this paper, we have identified a new flare star 2E 2206.6+4517 using extensive ROSAT X-ray observations and optical spectra. The object has been classifed as a M3.0Ve star.  Three large X-ray flare events are seen with light curves characteristic of outbursts detected in other flare stars. Two flares with almost complete X-ray coverage show a variation in strength and timescales.  A flare detected with the PSPC had a peak luminosity L$_{X}$=1.1$\\times10^{30}$ erg s$^{-1}$, an e-folding rise time of $\\approx$2.2 hours and a decay time of $\\approx$7 hours.  We interpret this long decay time (one of the longest ever observed) as possible evidence for continual heating as similar to recent observations and analysis of AD Leo and EV Lac (Favata et al., 2000; Schmitt, 1994).  An observation with the HRI detected a flare with a higher peak luminosity of L$_{X}$=2.9$\\times10^{30}$ erg s$^{-1}$, an e-folding rise time of $\\approx$15 minutes, and a decay time of $\\approx$1.2 hours.  We used a statistical analysis of the X-ray light curves to measure the quiescent X-ray level, the (minimum) fraction of counts in flares and (minimum) fraction of time that the star exhibits flare or microflare activity.  The PSPC data showed a quiescent luminosity of $L_X = 4.6\\pm 0.2 \\times 10^{28}$ erg s$^{-1}$, with $\\geq$68$\\pm$2\\% of counts coming from flares, and significant flare contribution to at least 47$\\pm$5\\% of the time observed. The HRI data, taken on average almost three years later, show a quiescent luminosity of $L_X = 5.4\\pm 0.4 \\times 10^{28}$ erg s$^{-1}$, with $\\geq$59$\\pm$2\\% of counts coming from flares, and significant flare contribution to at least 41$\\pm$8\\% of the time observed.  We obtained two optical spectra that show strong and variable emission lines of the hydrogen Balmer series and neutral helium which is further evidence for flare activity.  With the large amount of X-ray coverage of the source 2E 2206.6+4517 in the ROSAT data archive, this object provides an excellent opportunity to facilitate in the understanding of the physical properties of these flare stars including the flare, quiescent and long term activity.  Further measurements of this object may be afforded by calibration observations of AR Lac by $\\textit{Chandra}$ or XMM-$\\textit{Newton}$."
    },
    "0011/astro-ph0011211_arXiv.txt": {
        "abstract": "The radio images from the Parkes-MIT-NRAO (PMN) Southern Sky Survey at 4850 MHz have revealed a number of previously unknown radio sources. One such source, G312.5-3.0 (PMN J1421-6415), has been observed using the multi-frequency capabilities of the Australia Telescope Compact Array (ATCA) at frequencies of 1380 MHz and 2378 MHz. Further observations of the source were made using the Molonglo Observatory Synthesis Telescope (MOST) at a frequency of 843 MHz. The source has an angular size of $18\\arcmin$ and has a distinct shell structure. We present the reduced multi-frequency observations of this source and provide a brief argument for its possible identification as a supernova remnant. ",
        "introduction": "Extensive radio surveys of the sky began in the 1950s when the number of identified sources was not sufficient to conduct significant statistical studies. The large-area continuum radio surveys have become the foundation for all other radio astronomy tests since they provide the necessary radio source catalogues and the location of objects with which astronomers can calibrate their instruments.  The Parkes-MIT-NRAO (PMN) survey of the southern sky was undertaken as a collaboration between the Parkes Radio Observatory of the Australia Telescope National Facility (ATNF), the Massachusetts Institute of Technology (MIT), and the National Radio Astronomy Observatory (NRAO) \\citep{gri93}. The primary aim of the project was to survey the whole southern sky at a frequency of 4850 MHz in order to provide a more extensive, high-quality catalogue of radio sources for the Southern Hemisphere. The secondary aim was to use this survey to complement the results obtained from the Northern Hemisphere 4850 MHz survey performed by \\citet*{con89} which was made using the NRAO 91 m telescope \\citep{wri94}. Previously, the only large-area, high-frequency radio survey in the same area was the 2700 MHz survey of Bolton and his collaborators which contains around 8200 sources \\citep{gri91}.  The Southern Survey $(-87.5\\degr<\\delta<-37\\degr)$ and the Tropical Survey $(-29\\degr<\\delta<-9.5\\degr)$ increased the number of known radio sources in the Southern and Tropical zones by a factor of about 5 over those known previously from the Parkes 2700 MHz and Molonglo 408 MHz surveys \\citep{gri93}. The catalogue of radio sources discovered in the Southern Survey contains 23,277 sources. However, it is important to realise that the PMN surveys primarily detected point sources and the reduction process was optimised for such objects. This resulted in the exclusion of many extended objects from the survey with the exception of stronger sources which had a sufficiently high signal-to-noise ratio \\citep{wri94}.  \\begin{figure} \\psfig{file=figure1.eps,width=7.8cm} \\caption{A contour map from the PMN survey showing the object of interest. The contour levels start at 10 mJy/Beam and increase in steps of 10 mJy/Beam.} \\end{figure}  It was from an image of the PMN survey that the chosen object of study was discovered. The object is at right ascension $14^{\\mathrm{h}}20^{\\mathrm{m}}$ and declination $-64.2\\degr$ (Galactic longitude $312.5\\degr$ and latitude $-3\\degr$) and has an approximate diameter of $18\\arcmin$. The object is shown towards the centre of Figure~1 which is a map taken from the PMN survey. The circular appearance of the object suggests it may be a supernova remnant (SNR) although it does not currently appear in SNR catalogues. The source has also been noted as a SNR candidate by \\citet{dun97}. The object seems to have no optically visible counterpart and no counterpart has been found in current ROSAT data. Shown in Figure 2 is an IRAS 12 micron image from the same region of sky as the object \\citep{mcg96}. This indicates that the object may have a strong counterpart in the infra-red, although dust in the local environment makes further conclusions based on the IR data difficult.  \\begin{figure} \\psfig{file=figure2.eps,width=7.8cm} \\caption{A contour map from the IRAS 12 micron survey showing the same region of sky. The contour levels start at 380 mJy/$\\sq\\arcmin$ and increase in steps of 8 mJy/$\\sq\\arcmin$.} \\end{figure} ",
        "conclusions": "Mosaic observations of an extended source from the PMN survey were carried out using the ATCA in early 1995 at frequencies of 1380 MHz and 2378~MHz. To complement these data, further observations of the source were obtained from the MOST at a frequency 843~MHz.  The morphology of the extended object shows a distinct shell structure which is optically thin. This is common to both SNRs and planetary nebulae. In the case of a SNR, a shell structure would indicate that the remnant was possibly produced by a Type I supernova resulting from the collapse of a white dwarf star. Interactions with the interstellar medium may have produced variations in the spherical nature of the remnant's structure.  If the source is a SNR, the lack of significant polarisation suggests that the intervening magnetic fields are disordered or that the SNR candidate is at a large distance.  An estimate of the source's distance was made using the $\\Sigma-D$ relation for Galactic shell remnants. This yielded a diameter of $\\approx 53 \\, \\mathrm{pc}$ and a distance of $\\approx 10 \\, \\mathrm{kpc}$. These results are consistent with the known parameters of SNRs and that the apparent lack of linear polarisation is likely due to the large distance. The distance estimate is of course dependant on the source being a SNR.  Since interferometers do not respond to the shortest spatial frequencies, it is difficult to compare total flux densities from ATCA observations with those from the Parkes Radio Telescope. The 1.38 GHz and 2.38 GHz results may be improved by using the Parkes telescope which contains the lower spacings. Additional spectral information may be obtained from further observations using the ATCA and the Parkes telescope at X-Band (8000--9200 MHz). It would also be useful to obtain X-ray spectra of the source to gain an estimate of the shock velocity and thus determine the age and expansion rate of the SNR candidate."
    },
    "0011/astro-ph0011357_arXiv.txt": {
        "abstract": "{ The high signal to noise and good point spread function of XMM have allowed the first detailed study of the interaction between the thermal and radio emitting plasma in the central regions of M87. We show that the X-ray emitting structure, previously seen by ROSAT, is thermal in nature and that the east and southwest extensions in M87's X-ray halo have a significantly lower temperature (k$T=$ 1.5 keV) than  the surrounding ambient medium (k$T=$ 2.3 keV). There is little or no evidence for non-thermal emission with an upper limit on the contribution of a  power  law component of spectral index flatter than 3 being less than  1$\\%$ of the flux in the region of the radio lobes. ",
        "introduction": "The giant elliptical galaxy M87 is located at the center of the  X-ray diffuse emission of the irregular Virgo Cluster (Fabricant et al. 1980). Its nearby position ($\\sim$ 17-20 Mpc; Freedman et al. 1994, Tammann \\& Federspiel 1997) and the XMM-EPIC sensitivity allow detailed measurements of the spectral structure of the X-ray emitting gas on scales down to $\\sim$1 kpc.  Recently, Harris et al. (2000) using ROSAT/HRI data  and the recent high resolution 90 cm map of Owen et al. (2000) find general, but not precise coincidence between some radio and X-ray emitting regions (features described as the eastern `ear' and the southern `cobra'). Earlier,  B\\\"ohringer et al. (1995) using ROSAT/PSPC  showed that the lack of detailed spatial correlation of the radio and X-ray emission and the apparent  thermal nature of the excess emission in the X-ray asymmetric structure  exclude the inverse Compton (IC) mechanism as the origin of most of the X-ray radiation.  The observed X-ray density and temperature distribution  of the large scale emission (Mushotzky \\& Szymkowiak 1988 and references therein)  suggests a radiative energy loss of the order of 10$^{43}$ ergs s$^{-1}$ which implies that the gas is cooling and it was well modeled, before the XMM data, by a classical  cooling-flow model (Stewart et al. 1984, Fabian et al. 1984). In the same region of space a pair of relativistic jets is emerging from the active nucleus of M87, ejecting  energy into this environment at an estimated rate of  10$^{44}$ ergs s$^{-1}$ in the form of relativistic plasma (e.g. Owen et al. 2000). How these two media interact and if this leads to a heating of the intra-cluster plasma is still a matter of debate.  Detailed analysis of the radio data  (Owen et al. 1999,2000) finds a competition between inflow of the hot cluster gas and violent outflow of energy from the inner region, but the relation between the radio jet and the outer halo is still not fully understood.  The  detailed morphology of the radio emission suggests the  ``buoyant bubbles'' interpretation, as  first proposed for radio lobes in galaxies by Gull and Northover (1973). Churazov et al. (2000) have recently modeled this process with specific application to M87 to explain both the radio lobes and the  surrounding  X-ray structures. In this model, the ambient gas is captured and uplifted by the relativistic gas; the bubble expands and transforms itself into a torus rising in the potential well,  and one expects to find  thermal  gas, originating in the central regions, in the cavity of the torus. This dynamical behavior  implies that the high surface brightness regions are from the uplifted gas and that this effect is increased by cooling-flow conditions.  XMM-Newton observations, which provide unprecedentedly good combined spectral and imaging capabilities, well matched to the spatial scale sizes of interest,  are of key importance in clarifying the nature of M87's X-ray emission.  In this paper, we focus our study on the asymmetric X-ray enhancements detected by XMM-Newton. The core and jet data, together with the characteristics of the underlying extended gas distribution are presented in  B\\\"ohringer et al. (2001) in this volume. High spectral resolution RGS data will be published in a later paper by Sakelliou et al. (2001).  The paper is organized as follows: we present the observations in section 2. Section 3 describes the spectro-imaging data analysis. Section 4 is dedicated to the scientific discussion and to the relationship between the X-ray and radio emission. The conclusions are in section 5. ",
        "conclusions": "XMM/Newton observations of the X-ray extended features of M87 have definitively established the thermal nature of the plasma, with the FeL complex dominating the emission. The temperature in these regions is significantly lower than the surrounding medium. This result is very robustly established thanks to the high sensitivity of the experiment. The model of Churazov et al.  seems to qualitatively explain our result for the eastern arm but fails to describe the misalignment between X-ray and radio lobes."
    },
    "0011/astro-ph0011343_arXiv.txt": {
        "abstract": "We present a new model for time evolution of fast magnetic reconnection in free space, which is characterized by self-similarity. Reconnection triggered by locally enhanced resistivity assumed at the center of the current sheet can self-similarly and unlimitedly evolve until external factors affect the evolution. The possibility and stability of this type of evolution are verified by numerical simulations in a very wide spatial dynamic range. Actual astrophysical reconnection in solar flares and geomagnetospheric substorms can be treated as an evolutionary process in free space, because the resultant scale is much larger than the initial scale. In spite of this fact, most of the previous numerical works focused on the evolutionary characters strongly affected by artificial boundary conditions on the simulation boundary. Our new model clarifies a realistic evolution for such cases. The characteristic structure around the diffusion region is quite similar to the Petschek model which is characterized by a pair of slow-mode shocks and the fast-mode rarefaction-dominated inflow. However, in the outer region, a vortex-like return flow driven by the fast-mode compression caused by the piston effect of the plasmoid takes place. The entire reconnection system expands self-similarly. ",
        "introduction": "\\label{sec:Int}  It is widely accepted that magnetic reconnection takes place very commonly as a very powerful energy converter in astrophysical plasma systems. However, there still remain so many open questions not only for the microscopic physics of the resistivity, but also for the macroscopic magnetohydrodynamical (MHD) structure.  One of the most important questions on magnetic reconnection is what condition determines the reconnection rate, by which we can estimate the speed of energy conversion. The reconnection rate is usually defined as the ratio of the Alfv\\'{e}n wave transit time scale across the system length to the energy conversion time scale. The reconnection rate mainly depends upon the magnetic Reynolds number $R_m$ which represents the ratio of the diffusion time scale of the magnetic field to the Alfv\\'{e}n wave transit time scale. We should note that the magnetic diffusivity of astrophysical plasma systems is very small ($R_m \\sim 10^{14}$ for the solar corona). We here review the previous models of magnetic reconnection from the viewpoint of the reconnection rate.  Models in which the magnetic energy is mainly converted by magnetic diffusion (e.g., the Sweet-Parker model [Sweet 1958, Parker 1963]) are called ``slow'' reconnection (Tajima \\& Shibata 1997), because these models crucially depend upon $R_m$ (i.e., the reconnection rate of the Sweet-Parker model is in proportion to ${R_m}^{-1/2}$). Hence, the reconnection rates of these models are very small, and these models cannot convert the magnetic energy in a short period (e.g., the typical duration of solar flares is $\\sim$ a few minutes-a few hours). These slow reconnection models cannot provide adequate reconnection rates for actual powerful energy conversion.  A reconnection model with a pair of slow-mode shocks is firstly presented by Petschek (1964). In this model, the magnetic energy is mainly converted at the slow-mode shocks. Therefore, the time scale of the energy conversion is determined by MHD wave propagation, and is much shorter than that of the simple diffusion model or of the Sweet-Parker model. The reconnection rate is almost independent of the magnetic Reynolds number ($\\propto \\log R_m \\sim {R_m}^0$). Hence, we call this quick magnetic energy converter ``fast reconnection''.  We should note, however, that the question of how we can attain the fast reconnection is still open. Concerning this point, there are two different models. One states that external boundary conditions should control the reconnection, so that the resistivity has no effective influence on the energy conversion (see, e.g., Petschek 1964, Priest \\& Forbes 1986 for theoretical studies, and Sato \\& Hayashi 1979 for numerical studies). The other has been proposed in a series of numerical works originated by Ugai \\& Tsuda (Ugai \\& Tsuda 1977, 1979, Tsuda \\& Ugai 1977, and recently, Ugai 1999). They put a localized finite resistivity in the current sheet assuming that such resistivity simulates the anomalous resistivity in plasmas. This resistivity acts as a trigger to start the magnetic reconnection. Once this resistivity is put, the reconnection starts and evolves self-consistently, and forms a fast reconnection system with the Petschek-type structure (a pair of slow-mode shocks is formed along the current sheet).  However, even in the previous numerical studies aimed at clarification of the time evolution of fast reconnection, the evolution could not be followed for a long time. This is mainly due to the finite size of the simulation box, hence the application of the results was limited to spatial scales typically, say, hundred times the spatial scale of the diffusion region.  The actual magnetic reconnection in astrophysical systems usually grows in a huge dynamic range in its spatial dimension. For example, the initial scale of the reconnection system can be defined by the initial current sheet thickness, but this is too small to be observed in typical solar flares. We do not have any convincing estimate of it, but if we estimate it to be of the order of the ion Larmor radius, it is so small ($\\sim 10^0$ m in the solar corona). Finally, the reconnection system develops on the scale of the order of the initial curvature radius of the magnetic filed lines ($\\sim 10^7$ m $\\sim$ 1.5\\% of the solar radius for typical solar flares). The dynamic range of the spatial scale is obviously huge ($\\sim 10^7$ for solar flares). For the geomagnetospheric substorms, their dynamic range of growth is also large ($\\sim 10^4$ for substorms). Such a very wide dynamic range of growth suggests that the evolution of the magnetic reconnection should be treated as a development in ``free space'', and boundary conditions do not affect the evolutionary process of magnetic reconnection, at least in the initial stage just after the onset of reconnection.  In the actual numerical studies, due to the restriction of computer memory, the dimension of the simulation box is limited.  For magnetic reconnection to be properly studied, the thickness of the current sheet must be sufficiently resolved by simulation mesh size. Usually, the thickness of the current sheet is much smaller than the size of the entire system. Hence, we are forced to cut a finite volume out from actual large-scale reconnection system for numerical studies.  Of course, we wish that the boundary conditions of this type of finite simulation box reproduces the evolution in unbounded space. However, we should note that, in actual simulations, the boundary of such a simulation box necessarily affects the evolution inside the box, even if we take so-called ``free'' boundary conditions. This fact has an adverse influence on the evolutionary process of the numerically simulated magnetic reconnection as follows. When the reconnection proceeds and physical signals propagating from the inner region cross the boundaries of the simulation box, the subsequent evolution is necessarily affected by the boundary conditions. This is because the information propagated through the boundary is completely lost and such an artificially cut-out simulation box will never receive the proper response of the outer regions. Additionally, numerical and unphysical signals which are emitted from the boundary may disturb the evolution. Such artificially affected evolution is obviously unnatural, and the resultant stationary state should be different from the actual reconnection occurring in free space.  In this paper, our interest will be focused on the evolutionary processes of magnetic reconnection in ``free space''. On the basis of numerical simulations, we propose a new model for time-dependent magnetic reconnection. In this model, an entire system of magnetic reconnection triggered by locally enhanced resistivity continues to evolve self-similarly and unlimitedly. We call our model the self-similar evolution model. Propagation of the fast-mode rarefaction wave emitted from the reconnection region plays crucial roles in this self-similar evolution model. This point has never been pointed out in the previous works. The launch of rarefaction waves induces a flow toward the current sheet, which leads to a characteristic inflow structure. We suggest that this characteristic structure is related to the observational results on solar flares (see section \\ref{sec:App}). Our reconnection system is accompanied by a pair of slow-mode shocks caused by the induced inflow just as in the Petschek model. Majority of the magnetic energy is converted at this slow shock. Hence, this is categorized as fast reconnection. The scenario for the evolutionary process is clarified in section \\ref{sec:Sce}, and properties of this new model are discussed in section \\ref{sec:S-D}. The results of numerical simulations performed in this paper are summarized in section \\ref{sec:Num}. On the basis of such numerical studies, we obtain an intuition to the model. This kind of intuition will stimulate our theoretical effort to find an analytic representation of the self-similar evolution model. ",
        "conclusions": "\\label{sec:S-D}  \\subsection{Summary} \\label{sec:Sum}  We have presented a new model for describing the evolutionary process of magnetic reconnection: the self-similar evolution model.  The possibility of this type of evolution has been demonstrated by numerical simulations. An outline of the evolutionary process has been given in section \\ref{sec:Sce}.  We should note that, in this new model, propagation of the fast-mode rarefaction wave (FRW) plays an important role in the self-similar expansion of the reconnection system. FRW propagates almost isotropically at a constant speed $V_{A0}$ (the Alfv\\'{e}n speed in the asymptotic region) for the case of low $\\beta$ plasmas ($\\beta \\ll 1$). The only characteristic scale of the system at a later stage (the self-similar evolution stage, see figure \\ref{fig:stage3f4}) is the radius of the circular wave front of FRW (FRWF) because the initial current sheet thickness $D$, which is a fixed characteristic scale, is negligible in comparison with the scale of FRWF in this stage. This fact suggests the possibility of a self-similar solution. In fact, the results of our numerical simulation (see figures \\ref{fig:series1f5}, \\ref{fig:series2f6}) clearly show that the evolution settles to a stationary state if they are expressed in the zoom-out coordinates (see section \\ref{sec:Sim}). This is equivalent to saying that such evolution is self-similar. Again note that the self-similar evolution is achieved only at the last stage of evolution ($t \\gg D/V_{A0}$) in which the fixed characteristic scale ($D$ or the size of the resistive region) is negligible as compared to the system size (the radius $V_{A0} t$ of FRWF).  Of course, at that stage, the shape and the size of the resistive region, which are artificially given in the model, no longer influences the evolution.  This new type of reconnection system represents ``fast'' reconnection. As FRW propagates to the uniform region (the asymptotic region $y \\gg D$), an inflow is induced by the decrease of the total pressure near the reconnection point. Figure \\ref{fig:det-magf14} clearly shows that a pair of slow shocks is formed by this induced flow toward the current sheet. These slow shocks efficiently convert magnetic energy to kinetic or thermal energy of plasma. Of course, the speed of this energy conversion is determined by the slow-mode wave speed, and is independent of the magnetic diffusion speed ($\\propto {R_m}^0$). Hence, this energy release is fast one, and we can call it fast reconnection.    \\subsection{Why self-similar?} \\label{sec:Why}  In general, self-similar solutions can exist in the systems having no fixed characteristic scale. Our two-dimensional current sheet system however has a fixed characteristic scale: the initial current sheet thickness $D$.  The typical scale of the entire system can be estimated from the radius of FRWF ($\\sim V_{A0} t$). When $t \\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}} D/V_{A0}$, the evolution is affected by initial conditions, but such effects gradually diminish as the system scale ($V_{A0} t$) exceeds the current sheet thickness ($D$). When $t \\gg D/V_{A0}$, the current sheet thickness is negligible comparing with the system scale. In this situation, the system no longer has any characteristic spatial scale than the radius $V_{A0} t$ of FRWF, and the evolution settles to a self-similar solution. In fact, we found self-similar solutions as discussed in section \\ref{sec:Num}. Also the size and the shape of the resistive region no longer influence the evolution, because the dimension of the resistive region is negligible as compared to that of the entire system. In fact, we can confirm that the evolution at the later stage is insensitive to the resistivity model (see figure \\ref{fig:case_c_pmf15}).  \\placefigure{fig:case_c_pmf15}  It is important to check whether such self-similar solutions are stable or not. We have verified by our numerical simulations that our self-similarly evolving solution is stable over $\\sim 10^2$ of the spatial dynamic range.  The existence of this self-similar evolution stimulates theoretical efforts to find an analytic solution to the problems of self-similar evolution. This will be discussed in a forthcoming paper.  Strictly speaking, the two-dimensional treatment (resistivity enhancement elongated in the $z$ direction) of this self-similar model may not be plausible for actual cases which will be triggered by a ``point'' resistivity enhancement. In this case, we must treat the evolution as a three-dimensional problem. Differences of such 3-D evolution from the 2-D self-similar evolution model will be treated as a future problem.  Biernat, Heyn and Semenov (1987) and Semenov et al. (1992) analytically studied the evolutionary process of magnetic reconnection in a similar situation to this paper. They treat the reconnection rate as a free parameter which can be varied arbitrarily. In this meaning, the case they treated is also categorized as spontaneous reconnection. Their analysis gives a general formalism of spontaneous time-varying reconnection. However, in their works, plasma of the inflow region is assumed to be incompressible for analytical convenience. This means that the sound speed is infinitely large even when the Alfv\\'{e}n speed is finite. This is equivalent to assuming that the inflow region is filled with extremely high $\\beta$ plasmas (note that $\\beta \\sim$ (sound speed)$^2$/ (Alfv\\'{e}n speed)$^2$). Needless to say, this assumption is unsuitable for most cases of astrophysical problems. Contrary to our case, fast-mode waves emitted from the central region instantly propagates to infinity. Although the proper spatial scale determined by the fast-mode wave propagation does not exist in their case,  there is a possibility that the structure formed by slow-mode wave can be self similar, because the propagation speed of the slow-mode wave is finite in their case. Our work is understood as an extension of their works to realistic low $\\beta$ astrophysical plasma systems.    \\subsection{Adequacy of resistivity model} \\label{sec:Ade}  In our simulation, localization and stationarity of the resistivity are essential assumptions. Since the resistivity cannot be derived from MHD equations, we must apply a certain model for resistivity. In the problem of magnetic reconnection in astrophysical plasmas, so-called anomalous resistivity plays a part in the magnetic diffusion. In order to make a realistic model, the physical process causing the anomalous resistivity is very important, but remains as an open question.  In addition, as discussed in section \\ref{sec:Int}, there is room for discussion concerning two different points of view about influences of the resistivity on the evolution of magnetic reconnection. Even in the case where resistivity does not play crucial roles, there is a prediction that the fast reconnection will never be attained in a system with uniformly distributed resistivity (Biskamp 1986, Scholer 1989, Yokoyama \\& Shibata 1994). On the other hand, in the case where resistivity plays crucial roles, the reason why the resistivity can be localized is still unclear. We know neither the origin of resistivity nor influence of resistivity on the reconnection system.  Our opinion on the resistivity model is that localization and stationarity of the resistivity are essential for the self-similar evolution but the evolution does not depend upon the size or shape of the resistive region nor the physical process causing the resistivity.  In order to justify our opinion, we must discuss our simulation results with other resistivity models. As an example for a different resistivity model (case B), we have performed the following case (Nitta 1988). We assume a circular resistive region $r \\leq D$  in the current sheet, where $r$ is the distance from the origin and $D$ is the initial thickness of the current sheet. In this resistive region, the resistivity is distributed as an exponentially decaying function of $r$. Simulation with this resistivity model also results in self-similar evolution of magnetic reconnection.  Another example (case C) we have tested has the same shape of the resistive region as the previous case (resistivity uniformly enhanced in the central square region), but the value of the resistivity is five times larger ($\\eta=0.5$ thus $R_m=4.9$) than in the previous case ($\\eta=0.1$ thus $R_m=24.5$) . The result (see figure \\ref{fig:case_c_pmf15}) is very similar to the original case (see figure \\ref{fig:series1f5}). The evolution does not depend crucially on the value of resistivity.  In conclusion, we cannot identify any essential difference among these three cases. Hence, we believe that any model with a localized resistivity  and an infinitely large system size will result in the self-similar evolution.  Magnetic reconnection needs some processes to keep electric field (reconnection electric filed) along the reconnection line in order to break up the frozen-in condition. In the diffusion region of the plasma system with extremely large magnetic Reynolds number like in astrophysical problems, macroscopic MHD approximation is no longer valid, and the reconnection electric filed should be caused by microscopic processes. Three resistivity models we have examined above are all based on the Ohmic term $\\eta J$ where $\\eta$ is the resistivity and $J$ is the current density in the diffusion region. In these models, the localized resistivity $\\eta$ is supposed to represent anomalous resistivity induced by anomalous collision between particles and waves as a result of micro-instabilities (e.g., the lower hybrid drift [LHD] instability [Huba et al. 1977, Shinohara et al.1998, Shinohara et al. 1999].  However, in actual plasma systems, the reconnection electric field might arise from other terms of the generalized Ohm's law (Biskamp 1997). Several collisionless processes to maintain the reconnection electric field have been proposed (e.g., the Hall current effect [Ma \\& Bhattacharjee 1996] or the electron inertia effect [Tanaka 1995], the whistler turbulence effect [Shay \\& Drake 1998]). Unfortunately, we do not have a definite strategy for including such reconnection electric field produced by above mentioned microscopic processes in the macroscopic MHD regime. Thus, we adopted a simplified resistivity form in our Ohm's law.  Of course, the actual resistivity in astrophysical plasma systems cannot be described by simple Ohm's law as adopted in our simulation, but our self-similar evolution only requires a contrivance to keep a finite electric field along the reconnection line. We should emphasize again that the self-similar evolution occurs in the last stage when the system scale becomes much larger than the initial current sheet thickness and the size of the diffusion region. Hence, such evolution can be described by macroscopic MHD, and does not depend upon whether the central resistivity is caused by collisional or collisionless processes.  In the case of evolution in free space, boundary conditions that are crucial to the ``driven reconnection'' models obviously do not play important roles. Hence the resistivity should control the evolution any way. However, we do not go into details of the nature of resistivity itself in this paper, and adopt a simple model of localized and stationary resistivity. The dependence of the self-similar evolution on various resistivity models will be discussed in our forthcoming paper in detail.    \\subsection{Comparison with previous stationary models} \\label{sec:Com}  There are several theoretical models for steady state magnetic reconnection. We compare our self-similar evolution model with these previous models. Our discussion is focused only on the fast reconnection, because the very quick energy conversion frequently observed in astrophysical phenomena suggests that fast reconnection should be considered as the responsible mechanism.  The Petschek model (Petschek 1964) is characterized by a pair of slow shocks and fast-mode rarefaction in the inflow region. This fast-mode rarefaction wave is produced at the central reconnection region. As a result of the fast-mode rarefaction, the gradient of the magnetic field strength near the neutral point decreases due to the bending of magnetic field lines (Vasyliunas 1975). This process limits the diffusion speed, and hence the reconnection rate.  The Sonnerup model (Sonnerup 1970) is developed from the Petschek model. This model is characterized by a pair of slow shocks and the hybrid of fast-mode and slow-mode rarefaction. The fast-mode rarefaction wave is produced at the central region as in the Petschek model, but the slow-mode rarefaction wave is injected from the boundary. Because of the hybrid nature of rarefaction, the gradient of the magnetic field strength near the neutral point does not decrease as in the Petschek model. Therefore the Sonnerup model can attain the maximum reconnection rate possible for magnetic energy converters (Priest \\& Forbes 1986).  Our self-similar evolution is never influenced by the boundary conditions, so the Sonnerup-type hybrid rarefaction does not take place at all in our case. The fast-mode rarefaction dominates in the vicinity of the diffusion region. As discussed in section \\ref{sec:Sim}, the central region of this self-similarly evolving system is of the Petschek-type. Therefore, the reconnection rate will be limited in a way similar to the original Petschek model. This point will be clarified in our following papers.  Far from the central region, there is a region in which the fast-mode compression takes place by the piston effect of the plasmoid. This fast-mode compression causes the vortex-like return flow (see figure \\ref{fig:det-totf13}). We can consider that this self-similar evolution includes the Petschek model as the inner solution. A combined feature of the original Petschek model and the vortex-like flow properly characterizes the evolutionary process, and the entire system unlimitedly expands self-similarly.    \\subsection{Applications to flares} \\label{sec:App}  This study has been motivated by the fact that actual reconnection has a very wide spatial dynamic range of evolution; in the case of geomagnetospheric substorms or the solar flares; the dynamic range is $10^4-10^7$. Therefore, the early stage of the evolution of magnetic reconnection can be approximated as the evolution into free space without any influence of boundary conditions. This approximation is justified as long as the spatial scale of the reconnection system is much smaller than that of the entire system (e.g., the radius of curvature of the initial magnetic field lines). In the case of large solar flares, the typical radius of curvature is $\\sim 10^{7-8}$ m and the typical Alfv\\'{e}n speed is $\\sim 10^6$ m s$^{-1}$. Hence we can predict from our self-similar model that during the interval $\\sim 10^{7-8}/10^6 \\sim 10^{1-2}$ s from the onset of magnetic reconnection, the evolution can be approximated by our self-similar evolution model, and the total power, integrated over the entire system, of the energy conversion increases in proportion to time. The authors expect that coming observations by Solar B will detect such phenomena.  Another property of our self-similar model is the characteristic structure of the inflow region. A vortex-like flow around the head of the plasmoid characterizes the evolutionary process (see figure \\ref{fig:det-totf13}). We may be able to verify our evolutionary model by detecting this kind of vortex flow.  It has recently been found that solar flares (and similar arcade formation events) are often associated with depletion of ambient plasma density, called dimming, during their main energy release phase (Tsuneta 1996, Sterling and Hudson 1997). Tsuneta (1996) attributed this to the result of reconnection inflow (see also Yokoyama and Shibata 1997), while Sterling and Hudson (1997)  interpreted this as being due to plasma ejection from the system. In the case of a large cusp-shaped flare (Tsuneta 1996), the time scale of the ``dimming'' is 10-20 minutes just at the onset of the flare, which is about 7-14 $t_A$ (where $t_A$ is the Alfv\\'{e}n transit time), and the ratio of depleted density to the initial density is $\\sim$ 0.2 during this period. Let us examine the possibility that the fast rarefaction wave associated with reconnection inflow can account for the observed \"dimming\" at the onset of flares. In figure \\ref{fig:dimmingf16}, predicted X-ray image (in relative measure) of the reconnection system is shown. This figure is drawn from the result of case A. Temperature range is assumed as $10^7-10^8$ K, which is plausible for solar flares. We can clearly find that a round dark region around the reconnection point expands as the fast mode rarefaction wave propagates isotropically. Our simulation results shown in figure \\ref{fig:dimmingf16} show that the ratio of the depleted density to the initial density in the rarefaction wave is of order of 0.1-0.2, consistent with observations (see figure \\ref{fig:profile_rhof8}).  In actual reconnection systems, the spatial scale (say $L$) of the region in which magnetic flux is piled up should be finite. This region is an energy reservoir of magnetic reconnection. Although the time scale of strict self-similar evolution in the bounded system ($L=finite$) is the Alfv\\'{e}n time ($t_A \\equiv L/V_{A0}$), our self-similar evolution  (especially near the reconnection point) may roughly hold more than the Alfv\\'{e}n time (up to a few to 10 $t_A$; see figure \\ref{fig:rec_ratef17}) even if the system is not in an exact free space. The solid line of figure \\ref{fig:rec_ratef17} shows the time variation of the electric field ($\\eta |\\mbox{\\boldmath$J$}|$) at the reconnection point ($x=y=0$) for the typical case (case A). The intensity of the electric field is a measure of the reconnection rate of the system. The dashed line (case D) shows the reconnection rate for the case where the lower half region ($y \\leq 3000$) is filled with uniformly magnetized plasma, while the upper half region ($y>3000$) is unmagnetized. In case D, the fast mode wave is reflected at the boundary at $y=3000 \\  (t=3000)$, and returns to the resistive region at $t=6000$. Hence, the physical difference between these cases (case A and D) will appear after $t=6000$. Figure \\ref{fig:rec_ratef17} shows that the difference is only a few percent, and moreover, we cannot find any significant difference between these two cases even after the reflected wave returns to the current sheet ($t \\geq 6000$).  Consequently, basic characteristics (density and time scale) of fast-mode rarefaction wave in a self-similar evolution of fast reconnection seem to account for observations of ``dimming'' associated with solar flares.  \\placefigure{fig:dimmingf16}  \\placefigure{fig:rec_ratef17}"
    },
    "0011/astro-ph0011496_arXiv.txt": {
        "abstract": "Primordial chemistry began, at the recombination epoch, when the adiabatic expansion caused the temperature of the radiation to fall below 4000K. The chemistry of the early Universe involves the elements hydrogen, its isotope deuterium, helium with its isotopic forms and lithium. In this contribution I will discuss the influence of the primordial molecules on the cosmological decoupling. \\\\ In the framework of the gravitational instability theory, each protostructure started as a tiny local overdensity. As long as these inhomogeneities are small, their evolution can be studied by the classical linear perturbation theory. Once the deviations become large, the linear theory is no more valid. we present the role played by these molecules on the transition between the linear regime and the non-linear regime, and show that the molecules can lead to a thermal change at the turn-around point between these two regimes. ",
        "introduction": "At early times the Universe was filled up with an extremely dense and hot gas. Due to the expansion it cooled below the binding energies of hydrogen, deuterium, helium, lithium which led to the formation of these nuclei (Sarkar 1996). After this nucleosynthesis period the recombination process is not instantaneous because the electrons, captured into different atomic energy levels, could not cascade instantaneously down to the ground state. Atoms reached the ground state either through the cosmological redshifting of the Lyman $\\alpha$ line photons or by the $2s-1s$ two photons process. Nevertheless the Universe expanded and cooled faster than recombination could completed, and small fraction of free electrons and protons remained. \\\\ The principles of calculations of the primordial recombination have been mentionned initially by Shklovskii (1967), Novikov \\& Zel'dovich (1967). Peebles (1968) was the first to present a theory in which the very complicated recombination process is reduced to simpler terms (see Puy \\& Signore 2000 for the historical description and references therein). \\\\ At the end of the recombination period, it is plausible to imagine that molecules could be formed. Temperature of matter and radiation as well as the density are not so high and we have possible collisional reactions between the species. However, in this cosmological context we have metal-free gas which does not allow the efficient reaction of adsorption on the surface of grains. \\\\ The literature on the chemistry in the post-recombination epoch has grown considerably in the recent years.  Many authors have developped studies of primordial chemistry in different contexts. For example Lepp \\& Shull (1984), Latter \\& Black (1991), Puy et al. (1993), Stancil et al. (1996), Galli \\& Palla (1998) for the chemical network; Palla, Galli \\& Silk (1995), Puy \\& Signore (1996, 1997, 1998a, 1998b), Abel et al. (1997) and Galli \\& Palla (1998) in the context of the formation of the first objects. \\\\ Chemistry of the early Universe is the gaseous chemistry of the hydrogen, helium, lithium and electrons species. The efficiencies of the molecular formation processes is controlled by collisions, matter temperature and temperature of the cosmic microwave background radiation. The complete chemical network consists of 90 reactions (see Puy et al. 1993 and more recently Galli \\& Palla 1998 and references therein). \\\\ Molecules could lead to radiative processes through the excitation of the rotational levels, and contribute to the thermal evolution of the medium. Puy et al. (1993) examined the thermal balance in the early Universe and showed that the molecular cooling and heating functions have a slight influence on the decrease of the temperature of the matter. Nevertheless molecules could play a more decisive role in the collapse of protoclouds, since molecular cooling can trigger thermal instability prior the gravitational instability, as mentionned Puy \\& Signore (1996, 1998a, 1998b). Protocloud evolute in three phases: \\begin{itemize} \\item linear evolution which approximatively follows expansion. \\item {\\it turn-around} epoch when the protocloud reaches its maximum value. \\item non-linear evolution of the collapse. \\end{itemize} At the turn-around point it is crucial to know the initial condition and particularly the temperature of the pre-collapse. In this contribution, I will recall the thermal function generated by the primordial molecules (Sect. 2), then I will show the influence on the thermal decoupling (Sect. 3). In Sect. 4, the role of the molecules on the turn-around point will be discussed and a brief outlook will be presented in Sect. 5. ",
        "conclusions": ""
    },
    "0011/astro-ph0011175_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gravitationally-lensed quasars are both valuable (individually and statistically) and rare. They are mostly discovered through time-consuming re-imaging of known quasars, but can also be found spectroscopically in galaxy redshift surveys. Q~2237+0305 (Huchra et al.\\ 1985) was found in this manner, but remains the only such discovery. Kochanek (1992) predicted that $\\sim 1$ in $10^5$ redshifts would yield a lens although this is somewhat pessimistic, as described below. It is also important that lenses discovered in galaxy surveys tend to have bright, low redshift lens galaxies. Such lenses are particularly useful, as shown by the successful measurement of the source-size and transverse velocity in the Q~2237+0305 system (e.g., Wambsganss et al.\\ 1990; Wyithe et al.\\ 1999). ",
        "conclusions": ""
    },
    "0011/astro-ph0011205_arXiv.txt": {
        "abstract": "We report the results of preliminary analysis of the {\\xmm} EPIC and RGS observations of the candidate black-hole binary \\lmcx3 between February and June 2000. The observations covered both the soft and the hard X-ray spectral states. The hard-state spectra were dominated by a power-law component with a photon index $\\Gamma \\approx1.9$. The soft-state spectra consisted of a thermal component with a multi-colour disk temperature $T_{\\rm in}$ of $\\sim 0.9$~keV and a power-law tail with $\\Gamma \\approx 2.7$. The model in which the X-rays from \\lmcx3 in the high-soft state are powered by a strong stellar wind from a massive companion is not supported by the small line-of-sight absorption ($n_{\\rm H} \\simlt  10^{21}$~cm$^{-2}$) deduced from the RGS data. The transition from the soft to the hard state appears to be a continuous process associated with the changes in the mass-transfer rate. ",
        "introduction": "\\lmcx3 (Leong et al.\\ 1971) is a persistent X-ray source in the Large Magellanic Cloud (LMC). It has an orbital period of 1.70~d, and a mass function $\\simeq 2.3 M_{\\odot}$ (Cowley et al.\\ 1983). The optical brightness of  $V \\sim 17$ indicates that the system has a massive companion (van Paradijs et al.\\ 1987) . The companion is often classified as a B3\\,V star (Cowley et al.\\ 1983), although a B5\\,IV spectral type has also been suggested (Soria et al.\\ 2001). The non-detection of eclipses in the X-ray curve implies that the orbital inclination of the system is \\simlt~$70^{\\circ}$ (Cowley et al.\\ 1983). The inferred mass of the compact star is \\simgt~$7 M_{\\odot}$ (Paczynski 1983), thus establishing that the system is a black-hole candidate (BHC).  Most BHCs show soft and hard X-ray spectral states. Their X-ray spectrum in the soft state generally consists of both a thermal and a power-law component. The thermal component can be fitted by a blackbody spectrum with a temperature $\\sim 1$~keV, and it is interpreted as thermal emission from the inner accretion disk. The power-law component is believed to be comptonised emission from a disk corona (Sunyaev \\& Titarchuk 1980) or from the high-speed infalling plasma near the black-hole event horizon (Titarchuk \\& Zannias 1998). The photon index $\\Gamma$ of the power law is $ \\approx 2.5$--4. In the hard state, the thermal component is insignificant. The spectrum is dominated by a flat ($1.5 \\simlt \\Gamma \\simlt 2$), extended power-law component.  \\lmcx3 is found in the soft state most of the time. LMC X-1, another high-mass BHC in the LMC, has been seen in the soft state only. In contrast, Cyg X-1, the high-mass BHC in our Galaxy, tends to be in the hard state for the majority of the time. Such preference for either the hard or the soft state is not evident in the low-mass transient BHCs (e.g.\\ GRO~J1655$-$40). It is still unclear how mass transfer occurs in \\lmcx3: whether via Roche-lobe overflow or capture of the stellar wind from the companion. Given the fact that a B3\\,V star is unable to drive a strong stellar wind or to fill its Roche lobe if it is in a binary system with a 1.7-d orbital period, it is difficult to explain the observed X-ray luminosity of the system.  In this letter, we report on the {\\xmm} observations of \\lmcx3 between 2000 February and June. The system was in an unusual hard state in April, and our observations span the transition from the soft state to the hard state and then back to a softer state. Making use of these new X-ray spectral data, we attempt to shed light on the question of mass transfer in the system (see also Soria et al.\\ 2001).  \\begin{table} \\begin{minipage}{80mm} \\caption{{\\xmm} EPIC Observation Log} \\label{mathmode} \\centering \\begin{tabular}{@{}lllll} \\hline Revolution/ID &  Instrument & Exposure \\footnote{For the PN small window mode, the live time was 71\\% of the exposure time listed here; for the MOS RFS mode the live time was 6.9\\% of the exposure time.} & Mode \\footnote{See {\\it XMM} Remote Proposal Submission Software Users' Manual} \\\\ \\hline \\hline && \\\\ Rev0028/201 &  PN & 08.1~ks    & small window \\\\[2pt] Rev0028/301 &  PN & 0.3~ks    & full PN      \\\\[2pt] &  MOS1 & 08.9~ks    & partial RFS  \\\\[2pt] Rev0030/501 & MOS1 & 05.9~ks    & partial W5   \\\\[2pt] Rev0041/101 &  PN & 0.7~ks    & full PN       \\\\[2pt] Rev0041/401 & MOS1 & 03.5~ks    & partial W2   \\\\[2pt] &  MOS1 & 03.6~ks    & partial W5   \\\\[2pt] &  MOS1 & 03.4~ks    & partial W4   \\\\[2pt] Rev0045/101 &  PN & 01.6~ks    & full PN      \\\\[2pt] Rev0045/301 & MOS1 & 06.6~ks    & partial W3   \\\\[2pt] Rev0066/101 &  PN & 26.6~ks \\footnote{Only the first 8.3~ks are useful} & small window \\\\[2pt] Rev0092/201 &  PN & 24.2~ks    & small window \\\\[2pt] &  MOS1 & 20.3~ks    & partial W4 \\\\[2pt] \\hline \\end{tabular} \\end{minipage} \\end{table}  \\begin{table} \\begin{minipage}{80mm} \\caption{{\\xmm} RGS Observation Log } \\label{mathmode} \\centering \\begin{tabular}{@{}lllll} \\hline Revolution/ID & Instrument \\footnote{The labels ``s\\#'' identifies multiple exposures during the same revolution; for more details see the Preferred Observation Sequence files of each revolution.} & Exposure & count rate \\footnote{In photons~s$^{-1}$} \\\\ \\hline \\hline && \\\\ Rev0028/201 & RGS2 (s1) & 17.6~ks & 4.79$\\pm$0.02 \\\\[2pt] Rev0030/501 & RGS2 (tot) & 12.1~ks & 4.52$\\pm$0.03 \\\\[2pt] Rev0045/101             & RGS1 (s1) & 06.4~ks & 3.59$\\pm$0.03 \\\\[2pt] & RGS2 (s1) & 06.4~ks & 3.56$\\pm$0.03 \\\\[2pt] Rev0045/201 & RGS1 (s1) & 06.7~ks & 3.62$\\pm$0.02 \\\\[2pt] & RGS2 (s1) & 06.8~ks & 3.57$\\pm$0.02 \\\\[2pt] Rev0045/301 & RGS1 (s1) & 03.0~ks & 3.66$\\pm$0.04 \\\\[2pt] & RGS1 (s3) & 03.0~ks & 3.57$\\pm$0.04 \\\\[2pt] & RGS1 (s5) & 03.0~ks & 3.40$\\pm$0.04 \\\\[2pt] & RGS2 (s2) & 00.3~ks & 3.63$\\pm$0.17 \\\\[2pt] & RGS2 (s4) & 03.0~ks & 3.26$\\pm$0.04 \\\\[2pt] & RGS2 (s6) & 00.3~ks & 3.18$\\pm$0.17 \\\\[2pt] Rev0066/101             & RGS1 (s7) & 44.3~ks & 0.0154$\\pm$0.0012 \\\\[2pt] & RGS2 (s8) & 44.3~ks & 0.0153$\\pm$0.0012 \\\\[2pt] Rev0092/101             & RGS1 (s2) & 51.9~ks & 1.32$\\pm$0.01 \\\\[2pt] & RGS2 (s2) & 51.9~ks & 1.53$\\pm$0.01 \\\\[2pt] Rev0092/201 & RGS1 (s4) & 25.6~ks & 1.56$\\pm$0.01 \\\\[2pt] & RGS2 (s5) & 25.6~ks & 1.34$\\pm$0.01 \\\\[2pt] \\hline \\end{tabular} \\end{minipage} \\end{table}   \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(8.8,13.5) \\put(0,9){\\special{psfile=asm_xmm.ps hoffset=-20 voffset=-15 hscale=37.5 vscale=40} } \\put(0,6){\\special{psfile=f_rgs_no.ps hoffset=-20 voffset=-15 hscale=37.5 vscale=40} } \\put(0,3){\\special{psfile=f_temp_no.ps hoffset=-20 voffset=-15 hscale=37.5 vscale=40} } \\put(0,0){\\special{psfile=f_gamma_no.ps hoffset=-20 voffset=-15 hscale=37.5 vscale=40} } \\end{picture} \\caption{ (First panel from the top): The {\\xmm} EPIC observations of \\lmcx3, marked by solid vertical lines, are shown and compared with the {\\xte}/ASM light curve. The three {\\xte}/PCA observations that confirmed the hard state of the system (Boyd et al.\\  2000) are marked by dotted vertical lines. (Second panel): The 0.35$-$2.5~keV light curve from the RGS observations. (Third panel): The evolution of the fit temperature of the thermal component in the EPIC-PN and MOS spectra. The open circles are the means of the fit disk-blackbody temperatures $T_{\\rm in}$ error-weighted for the observations during the same revolution. (Fourth panel): The evolution of the photon index $\\Gamma$ of the power-law tail. The data points (open circles) are error-weighted means. The open squares are the photon indices obtained by Boyd et al.\\ (2000) from the three {\\xte}/PCA observations on May 5.76, 10.01 and 13.94~UT; a model with a single power law was used.  } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "\\subsection{Soft and hard states}  Previous {\\xte}/PCA observations (Wilms et al.\\ 2000) have shown that when the source was bright (with {\\xte}/ASM count rates $\\simgt 3$~ct~s$^{-1}$), the disk blackbody component was prominent, and its fit temperature $T_{\\rm in} \\approx 1.0$~keV. There was also a power-law tail with $\\Gamma \\approx 4$ in the spectra. The temperature $T_{\\rm in}$ appeared to decrease when the X-ray luminosity decreased, while the normalisation parameter $A_{\\rm disk}$ remained approximately constant. When the {\\xte}/ASM count rate dropped below $\\approx 0.6$~ct~s$^{-1}$, $T_{\\rm in}$ was reduced to $\\approx 0.7$~keV. The transition was accompanied by the hardening of the power-law component. The photon index $\\Gamma$ became $\\sim 2.0$--$3.0$. When the {\\xte}/ASM count rate decreased below the 0.3-ct~s$^{-1}$ level, the disk blackbody component was not detected, and the spectrum was a power law with $\\Gamma = 1.8$.  The soft-to-hard transition was also seen in our data. The spectra obtained in the Rev0028--Rev0045 observations show a disk blackbody component with $T_{\\rm in} \\sim 0.6$--$1.0$, and a power-law component with $\\Gamma \\approx 2.5$. The {\\xte}/ASM count rate was below $2$~ct~s$^{-1}$ during the observations. The spectral properties are similar to those observed previously when the system had similar {\\xte}/ASM count rates.  The spectrum obtained in the Rev0066/101 observations is dominated by a power law with $\\Gamma = 1.9 \\pm 0.1$. The photon index is consistent with that observed in the previous hard states (Wilms et al.\\ 2000), while a lower value of $1.60\\pm0.04$ was obtained by Boyd et al.\\ (2000) on 2000 May 7. The $\\chi^2$ of the fit which includes a thermal disk black-body component (best fit $T_{\\rm in} = 0.14$~keV) is 58.486 for 52 degrees of freedom. If the thermal component is not included, then we obtain $\\chi^2 = 58.493$ for 54 degrees of freedom. The thermal component is therefore insignificant.  The system was at the transition from the hard state back to the soft state during the Rev0092 observations (see Figure~1). The power-law component had steepened, with $\\Gamma = 2.05 \\pm 0.02$ (for the PN data). The disk blackbody component reappeared, with a temperature $T_{\\rm in}\\approx 0.2$~keV, significantly lower than that before Rev0066.  In summary, between 2000 February and June \\lmcx3 underwent a transition from a soft to a hard state, and then in the process of returning to the soft state. During the soft-to-hard transition the fit temperature of the disk blackbody component decreased, and the power law component became harder. As the system started to return to the soft state, the disk blackbody component became more prominent. The power law appeared to be steeper than that obtained from the RXTE/PCA observations near the middle of the faint state, yet the disk blackbody temperature was still well below the values of the previous soft state.   \\subsection{Mode of mass transfer}  The RGS data constrain the value of $n_{\\rm H}$ within the system to be $\\simlt 10^{21}$~cm$^{-2}$, in contrast to the larger intrinsic column density expected for a companion with a strong stellar wind. The non-detection of obvious emission lines in the RGS spectra also indicates the absence of wind matter ejected in previous epochs (cf. the P Cygni lines seen in Cir X-1, Brandt \\& Schulz 2000). Thus, the {\\xmm} spectral data do not support the mass-transfer scenario for \\lmcx3 in which the black hole accretes matter mainly from a strong stellar wind from a massive companion. The high luminosity observed in the X-ray bands therefore requires the companion to overflow its Roche lobe.  A similar conclusion is also obtained independently from an analysis of the optical/UV properties of the system (Soria et al.\\ 2001). The OM data obtained in the Rev0066 suggest that the companion is a B5 subgiant instead of a B3 main-sequence star. Roughly 3\\% of the X-rays would be intercepted by the companion, so that the rate at which energy is deposited into its atmosphere can be $\\simgt 5\\times 10^{36}$~erg s$^{-1}$ in the soft state. This rate is larger than the intrinsic luminosity of the companion. If the companion in \\lmcx3 is indeed a subgiant star its tenuous envelope is susceptible to irradiation heating. The soft-to-hard transitions seen in the {\\xte} observations in 1997/1998 and the {\\xmm} observations in 2000 may be caused by variations in the rate of mass overflow from the Roche lobe of the subgiant companion.  Wilms et al.\\ (2000)'s interpretation of the decreases in the RXTE/ASM count rate as evidence for transitions from soft to hard state is consistent with our data. We further propose that the decrease in the X-ray luminosity is caused by the decrease in the fraction of the Roche lobe filled by the companion star. When the companion is detached from its critical Roche surface, mass transfer will be dominated by a focused wind.  It is worth noting that the three known high-mass BHCs are all persistent X-ray sources which show preferential X-ray spectral states. While Cyg X-1 tends to be in the hard state, LMC X-1 and \\lmcx3 are more often found in the soft state. Recent studies (e.g.\\ Igumenshchev, Illarionov \\& Abramowicz 1999; Beloborodov \\& Illarionov 2000) have shown that accretion of matter with low angular momentum will give rise to hard X-rays instead of soft X-rays. The relative angular momentum of the accreting matter is smaller for wind accretion than for Roche-lobe overflow. Cyg X-1 has a 33-$M_{\\odot}$ O-type companion (Giles \\& Bolton 1986), which has a strong stellar wind. The companion stars in LMC X-1 and \\lmcx3 are less massive ($\\simlt 10$~M$_\\odot$) B stars, whose stellar wind is much weaker. Therefore, we suggest a unified scenario which relates the mode of mass transfer (wind or Roche-lobe overflow) to the spectral state preferentially observed in these three high-mass BHC binaries (hard or soft respectively). The subgiant companion of \\lmcx3 may occasionally underfill its Roche lobe because of feedback irradiative processes or instabilities in its envelope. This leads to the residual accretion of the (focused) wind matter which has relatively low angular momentum."
    },
    "0011/astro-ph0011519_arXiv.txt": {
        "abstract": "General relativistic spectra from accretion disks around rotating neutron stars in the appropriate space-time geometry for several different equation of state, spin rates and mass of the compact object have been computed. The analysis involves the computation of the relativistically corrected radial temperature profiles and the effect of Doppler and gravitational red-shifts on the spectra. Light bending effects have been omitted for simplicity. The relativistic spectrum is compared with the Newtonian one and it is shown that the difference between the two is primarily due to the different radial temperature profile for the relativistic and Newtonian disk solutions.  To facilitate direct comparison with observations, a simple empirical function has been presented which describes the numerically computed relativistic spectra well. This empirical function (which has three parameters including normalization) also describes the Newtonian spectrum adequately. Thus the function can in principle be used to distinguish between the two. In particular, the best fit value of one of the parameters ($\\beta$-parameter) $\\approx 0.4$ for the Newtonian case, while it ranges from $0.1$ to $0.35$ for relativistic case depending upon the inclination angle, EOS, spin rate and mass of the neutron star. Constraining this parameter by fits to future observational data of X-ray binaries will indicate the effect of strong gravity in the observed spectrum. ",
        "introduction": "\\label{sec: I} X-ray binaries are believed to harbor black holes or weakly magnetized neutron stars with an accretion disk. The X-ray emission arises from the hot ($\\approx 10^7$ K ) innermost region of the disk. In the case of a neutron star there will emission, in addition, from a boundary layer between the accretion disk and neutron star surface. Since the observed emission arises from regions close to compact object, these sources are possible candidates for studying strong field gravity.  In the standard theory (Shakura \\& Sunyaev 1973), the accretion disk is assumed to be an optically thick Newtonian one. In this model, the local emergent flux (assumed to be a blackbody) is equated to the energy dissipation at a particular radial point in the disk.  The observed spectrum is then a sum of black body components arising from different radial positions in the disk. General relativistic effects modify this Newtonian spectrum in two separate ways. Firstly, the local energy dissipation at a radial point is different from the Newtonian disk, giving rise to a modified temperature profile. Secondly the observed spectrum is no longer a sum of local spectra because of effects like Doppler broadening, gravitational red-shifts and light bending. Modified spectra, incorporating these effects, but with different approximations have been computed by several authors (e.g. Novikov \\& Thorne 1973; Asaoka 1989) for accretion disks around rotating (Kerr) black holes. These computations confirm the expected result, that the relativistic spectral shape differs from the Newtonian one by around 10\\%. Thus, for comparison with observed data with systematic and statistical errors larger than 10\\%, the Newtonian approximation is adequate. Ebisawa, Mitsuda and Hanawa (1991) showed that for typical data from the Ginga satellite, the relativistic spectrum cannot be differentiated from the Newtonian disk spectrum. They also found that the relativistic spectrum is similar in shape (at the sensitivity level of Ginga ) to the Comptonized model spectrum. Although, Ginga was not sensitive enough to distinguish between the different spectra, better estimates of fit parameters like accretion rate and mass of the compact object were obtained when the data was compared to relativistic spectra rather than the standard Newtonian one.  The present and next generation of satellites (e.g. ASCA, RXTE, Chandra, XMM, Constellation-X) with their higher sensitivity and/or larger effective area than Ginga are expected to differentiate between relativistic and Newtonian spectra from low mass X-ray binaries (LMXB) and black hole systems. However, as pointed out by Ebisawa, Mitsuda and Hanawa (1991), the presence of additional components (e.g. boundary layer emission from the neutron star surface) and smearing effects due to Comptonization may make the detection ambiguous. Nevertheless, the detection of strong gravity effects on the spectra from these sources, will be limited by the accuracy of theoretical modeling of accretion disk spectra rather than limitations on the quality of the observed data. Thus, it is timely to develop accurate relativistically corrected spectra for comparison with present and future observations. Apart from the importance of detecting strong gravity effects in the spectra of these sources, such an analysis may also shed light on the geometry and dynamics of innermost regions of accretion disks.  Novikov \\& Thorne (1973) and Page \\& Thorne (1974) computed the spectra of accretion disks around rotating (Kerr) black holes.  This formalism when directly applied to rotating neutron stars provide only a first order estimate: the absence of an internal solution in the case of Kerr geometry, makes it difficult to obtain, in a straightforward manner, the coupling between the mass and the angular momentum of the central accretor. This coupling depends on one hand on the equation of state of neutron star matter and on the other on the proper treatment of rotation within general relativity. Equilibrium configurations of rapidly rotating neutron stars for realistic equations of state have been computed by several authors (Bonazzola \\& Schneider 1974; Friedman, Ipser \\& Parker 1986; Cook, Shapiro \\& Teukolsky 1994; Stergioulas \\& Friedman 1994; Salgado \\etal 1994a,b; Datta, Thampan \\& Bombaci 1998). One crucial feature in all these calculations is that the space--time geometry is obtained by numerically and self--consistently solving the Einstein equations and the equations for hydrodynamic equilibrium for a general axisymmetric metric. With the aim of modeling spectra of LMXBs, we attempt, in this paper, to compute the spectrum of accretion disks around rotating neutron stars within such a space--time geometry. This is particularly important since LMXBs are old (population I) systems and the central accretor in these systems are expected to have large rotation rates (Bhattacharya \\& van den Heuvel 1993 and references therein)  Computation of the spectra is numerically time consuming and hence direct fitting to the observed data is impractical. For the sake of ease in modeling, we also present in this paper, a simple empirical analytical expression that describes the numerically computed spectra. As shown later, the same expression (which has three parameters including normalization) can also describe the Newtonian spectra. In particular, the value of one of the parameters (called $\\beta$ parameter here) determines whether the spectrum is relativistically corrected or not. This will facilitate comparison with observational data since only this $\\beta$-parameter has to be constrained to indicate the effect of strong gravity in the observed spectrum.  The next section describes the method used to compute the spectra. In \\S 3 the results of the computation and the empirical fits are shown. \\S 4 is devoted to discussion and summary. ",
        "conclusions": "\\label{sec: 4}  In this paper, we have computed relativistic spectra from accretion disks around rotating neutron stars for the appropriate space-time geometry. Several different EOS, spin rates and mass of the compact object have been considered.  The Doppler and gravitational effects on the spectra have been taken into account, while light bending effects have been omitted for simplicity. The spectrum differs from the Newtonian one, with the main difference being due to the different radial temperature profile for the relativistic and Newtonian disk solutions.  A simple empirical function has been presented which describes the numerically computed relativistic spectra well. This will facilitate direct comparison with observations. The empirical function (eq.~[7]) has three parameters including normalization. Another important advantage of this function is that it also describes the Newtonian spectrum adequately, and the value of one of the parameters ($\\beta$-parameter) distinguishes between the two. In particular, the best fit $\\beta$-parameter $\\approx 0.4$ for the Newtonian case, while it ranges from $0.1$ to $0.35$ for relativistic case depending upon the inclination angle, EOS, spin rate and mass of the neutron stars.  In principle, for sufficiently high quality data, the effects of strong gravity on the disk spectrum can be detected using this empirical function as a fitting routine and constraining the $b$-parameter. However, it must be emphasized that there are several reasons why this may not be possible. There could be systems which have additional components in the X-ray spectra; for example boundary layer emission from the neutron star surface. Uncertainties in modeling these additional components may lead to a wider range in the best fit $\\beta$-parameter. Thus accurate spectra of the boundary layer (with relativistic corrections) is also needed for modeling these systems . Moreover, X-rays could be emitted from hotter regions (e.g an innermost hot disk or a corona) giving rise to a Comptonized spectra instead of the sum of local emission assumed here. In this case, the empirical fit will probably not describe the observed data well. It has been assumed here that the color factor is independent of radius. Shimura and Takahara (1995) have shown from numerical computation that this could be the case for an accretion disk in a Schwarzschild metric. Apart from the fact that this was done for Schwarzschild metric, their numerical computation depends on the vertical structure of the disk which in turn depends on the unknown viscosity mechanism in the disk. If the color factor has a radial dependence, the spectral shape might change which may be confused to be a relativistic effect.  Despite these caveats the method described in this paper will be a step forward in the detection of strong gravity effects in the spectra of X-ray binaries. Future comparison with high quality observational data, will highlight the theoretical requirements that have to be met, before concrete evidence for strong gravity are detected in these systems and the enigmatic region around compact objects is probed."
    },
    "0011/astro-ph0011033_arXiv.txt": {
        "abstract": "We report here the flux densities of the evolving radio source \\SN\\ at \\freq\\ measured from observations made with the Molonglo Observatory Synthesis Telescope between 1994 September and 2000 May. The radio light curve shows that the rate of increase of the flux density jumped markedly around days 2800-3000 (i.e.\\ in the first half of 1995), and that since then the radio evolution has been remarkably well fitted by a simple linear increase of $(62.7\\pm 0.5)\\, \\mu{\\rm Jy \\, day^{-1}}$. We discuss in detail the relationship between the radio light curve and the recent brightening of the system at optical wavelengths. ",
        "introduction": "Supernova 1987A (\\SN ) in the Large Magellanic Cloud is the nearest supernova explosion detected in over 300 years. Although not particularly bright, its proximity has facilitated studies at a level of detail unprecedented in supernova research, and rivalled only by the bright SN1993J.  We report here on observations of the radio emission from \\SN\\ at \\freq\\ using the Molonglo Observatory Synthesis Telescope (MOST). Our results cover the period 1994 September -- 2000 May, extending previously published data for 1987 February -- 1994 September (Ball et al.\\ 1995). Together these data sets from the MOST comprise a unique time series, sampled at intervals of 2 to 6 weeks.  Radio emission from SN1987A was first detected with the MOST on 1987 February 25.23 UT (Turtle et al.\\ 1987), just two days after the arrival of neutrinos on February 23.32 (Bionta et al.\\ 1987). After rising quickly to a peak flux density of $\\sim 130$\\,mJy at \\freq\\ the emission decayed rapidly and was undetectable at radio frequencies after 1987 September. On 1990 July 5 MOST observations indicated the reappearance of radio emission from \\SN\\ (Ball et al.\\ 1995), an event confirmed about a month later by observations at 4.8\\,GHz with the Australia Telescope Compact Array (Staveley-Smith et al.\\ 1992). These observations were the first indications that the expanding shock from the supernova explosion had encountered a dramatic change in the circumstellar material (CSM) it was passing through, an interpretation first suggested on the basis of modelling of the radio emission (Ball \\& Kirk 1992). Confirmation of this scenario followed from measurement of the low expansion speed of the radio source (Staveley-Smith et al.\\ 1993; Gaensler at al.\\ 1997) which required a dramatic deceleration of the supernova shock at or prior to the onset of the second phase of radio emission. More detailed models for the radio emission (Duffy, Ball \\& Kirk 1995) and for the hydrodynamics of the interaction of the supernova ejecta with the CSM (e.g.\\ Chevalier \\& Dwarkadas 1995; Borkowski, Blondin \\& McCray 1997) all attribute the radio emission to the interaction of the SN shock with a density jump in the CSM.  The implied radius of the circumstellar density jump responsible for the reappearance of the radio emission is smaller than that of the dense ring surrounding \\SN\\ (Jakobsen et al.\\ 1991; Plait et al.\\ 1995; Sonneborn et al.\\ 1997). Spatially-resolved spectra of \\SN\\ from the STIS instrument on the Hubble Space Telescope indicated that on 1997 April 26 (day number 3714) the reverse supernova shock radius was $\\sim 80 \\%$ of that of the ring (Sonneborn et al.\\ 1998). HST images from the WFPC2 instrument show that line emission from a local `knot' on the inner edge of the ring brightened by tens of percent between early-1996 and mid-1997 (days 3200--3750; Garnavich, Kirshner \\& Challis 1997). There are now clear indications that similar brightenings are taking place over a large fraction of the ring (Bouchet et al.\\ 2000; Lawrence \\& Crotts 2000; Maran, Pun \\& Sonneborn 2000; Garnavich, Kirshner \\& Challis 2000), suggesting that the interaction of the supernova shock and the inner edge of the ring is underway. Further analysis of archival WFPC2 images shows that the optical signatures of the interaction between material from the supernova explosion and the inner edges of the ring have been detectable since early 1995 (Lawrence et al.\\ 2000).  The radio data we report here show that the emission from \\SN\\ at \\freq\\ has continued to increase monotonically from 1990 July until May 2000. We present 106 new flux density determinations for the radio emission from \\SN\\ at \\freq\\ between 1994 September and 2000 May. In \\S 2 we briefly describe the observations themselves. In \\S 3 (and Appendix A) we explain the approach used to reduce the observations of this time-varying source. The reduced data are described in \\S 4 and tabulated in Appendix B. In Appendix C we discuss an alternative, more direct data reduction method, and compare the results with those given in Appendix B. In \\S 5 we discuss the implications of the observed evolution of the 843 MHz radio flux density from \\SN. In particular we investigate the evidence for features in the 843 MHz radio light curve that may be related to the dramatic changes in the optical line emission from this source that herald the collision of the supernova shock with the ring. ",
        "conclusions": "\\subsection{Up to 1995}  The reappearance of radio emission from \\SN\\ in 1990 July was first revealed by MOST observations at \\freq , and is now attributed to the encounter of the expanding supernova shock with a circumstellar density jump well inside the circumstellar ring (Ball \\& Kirk 1992 -- hereafter BK92; Chevalier 1992; Duffy, Ball \\& Kirk 1995 -- hereafter DBK95). This interpretation is supported by the analysis of images of the radio emission from \\SN\\ made at $8.6\\,\\rm GHz$ using the Australia Telescope Compact Array (Staveley-Smith et al.\\ 1993; Gaensler et al.\\ 1997). In particular, model fits to super-resolved images of the radio emission suggest that the source of the radio emission was expanding at only $\\sim 2800\\,\\rm km\\,s^{-1}$ in the period 1992--1995, much lower than the average speed of $35,000\\,\\rm km\\,s^{-1}$ required between the time of explosion and the first imaging observations in 1992 (Gaensler et al.\\ 1997). While this expansion speed is not a direct measurement of the speed of the supernova shock, it clearly indicates that the shock had undergone a dramatic deceleration from its initial speed of greater than $50,000\\,\\rm km\\,s^{-1}$ (Hanuschik \\& Dachs 1987), as would occur if the shock encountered a significant jump in the density of circumstellar material (Chevalier \\& Dwarkadas 1995; Borkowski, Blondin \\& McCray 1997) associated with a dense H{\\sc ii} region inside the circumstellar ring. The interpretation is also consistent with the observed soft X-ray emission from \\SN\\ (Beuermann, Brandt \\& Pietsch 1994; Gorenstein, Hughes \\& Tucker 1994; Hasinger, Aschenbach \\& Tr\\\"umper 1996), though the X-ray evolution is relatively poorly sampled.  The only attempts to model the evolving second phase of radio emission from \\SN\\ are those of BK92 and DBK95. Based on diffusive acceleration of electrons at an evolving supernova shock, including modification of the shock hydrodynamics by the pressure of accelerated protons, their model fits very well from switch-on to around day 2600 (as shown in Figure \\ref{light}). BK92 and DBK95 proposed that the switch on and subsequent rise of the radio emission resulted from the supernova shock encounter with a significant jump in the density of the circumstellar material. Chevalier \\& Dwarkadas (1995) placed this suggestion on a firmer footing, indicating that a very large density jump was expected at the inner boundary of a dense H{\\sc ii} region surrounding the progenitor. At times later than those considered by DBK95 (i.e. after about day 2600 or May 1995) the model flux density reaches a plateau and then decreases slowly -- contrary to the observations. The DBK95 model has many limitations (such as assuming spherical symmetry) which rapidly become less tenable at such epochs. No detailed model that combines a plausible distribution of the circumstellar material with a calculation of the supernova shock hydrodynamics and which treats the particle acceleration self consistently has yet been developed.  \\subsection{1995 on}  The optical brightening of the inner edge of the circumstellar ring that has been observed from mid-1996 on (Pun et al.\\ 1997; Garnavich, Kirshner \\& Challis 1997, 1999 \\& 2000; Bouchet et al.\\ 2000; Lawrence \\& Crotts 2000; Maran, Pun \\& Sonneborn 2000) clearly indicates that the interaction of the supernova shock and the ring is underway. H$\\alpha$ images show that an isolated spot just inside the main ring brightened by roughly 50\\% between 1994 February (day $\\sim 2500$) and 1997 July (day $\\sim 3750$) and broad band (R) images indicate a brightening by $\\sim 33$\\%, occurring between 1996 February (day $\\sim 3300$) and 1997 July (Garnavich, Kirshner \\& Challis 1997). More recent reports of R-band observations state that the brightening over the four years prior to 1999 February are well fitted by a steady increase in brightness, suggesting that the optical signature of the interaction may have begun as early as 1995 February (day $\\sim 2900$; Garnavich, Kirshner \\& Challis 1999). This has been confirmed by the analysis of archival HST data by Lawrence et al.\\ (2000).  It is possible that the jump in the rate of increase of the radio emission around days 2800 -- 3000 and the subsequent long-lived linear rise is associated with the impact of the supernova shock with the ring material. If not, the lack of features in the \\freq\\ light curve between days 2800 and 4820 implies that there is no significant radio emission from the impact sites evident in the optical images. Unfortunately, the highest resolution presently available at radio wavelengths with sufficient sensitivity is $\\sim 0.4$ arcsec at 9 GHz (Gaensler et al.\\ 1997) so direct comparison of radio and optical images is not definitive in this regard.  \\subsubsection{Timing and the shock radius}  The optical `hot spots' which were the first indications of the shock-ring interaction have been reported as being at position angle $\\sim 30^\\circ$ and $0.55\\;$arcsec from the supernova ($87\\pm 4$\\% of the ring radius; Garnavich, Kirshner \\& Challis 1997) and at p.a.\\ $103^\\circ\\pm5^\\circ$ and $0.6-0.65\\;$arcsec (around 80\\% of the ring radius; Bouchet et al.\\ 2000; Lawrence \\& Crotts 2000). Lawrence et al.\\ (2000) have shown that there was detectable optical emission from the first hot spot (near p.a.\\ $30^\\circ$) since at least March 1995 (day $\\sim 2950$). This first optical hot spot is not coincident with either radio lobe seen in the $9\\;$GHz images of Gaensler et al.\\ (1997). The brightening first seen around p.a.\\ $100^\\circ$ has now been shown to extend over position angles from $90^\\circ$ to $130^\\circ$ (Maran, Pun \\& Sonneborn 2000), which is roughly coincident with the brighter of the two radio `lobes' or `hot spots' seen in the $9\\;$GHz images of Gaensler et al.\\ (1997) -- and with the major axis of the ring. Lawrence et al.\\ (2000) have shown that the emission from this region began no later than 1999 January (day $\\sim 4350$) and probably as early as 1998 February (day $\\sim 4000$).  If the first optical hot spot is in the ring plane, its projected radial distance of $\\sim 0.55\\;$arcsec and p.a.\\ of $30^\\circ$ imply a true radial distance from the progenitor of $\\sim 0.7\\;$arcsec. The appearance of this emission in early 1995 (around day 3000) therefore suggests a shock radius which is consistent with or perhaps slightly larger than that inferred from model fits of a spherical source to the radio images (Gaensler et al.\\ 1997, Figure 10), given the significant uncertainties in both estimates. The near coincidence of the second more extended region of optical emission with the major axis of the ring implies that if the radiating material is in the ring plane, its true radial distance from the progenitor is the same as its projected radial distance in the plane of the sky. The appearance of this optical emission at a radial distance of $\\sim 0.65\\;$arcsec around the beginning of 1998 is again consistent with the radius of model fits to the radio images. Furthermore, the very similar radial distances to the first and second optical hot spots, and the significant delay between their appearance, are broadly consistent with the low expansion speed inferred from the fits to the radio images (Gaensler et al.\\ 1997).  \\subsubsection{Acceleration times}  The acceleration time required for electrons to attain sufficient energy to produce radio frequency synchrotron emission is an important physical parameter. The appearance of optical hot spots suggest that the shock impact with the ring began at least as early as March 1995 (day $\\sim 2950$; Lawrence et al.\\ 2000). If the change in the evolution of the \\freq\\ radio emission around day 3000 is related to the ring impact, the implied acceleration time is less than a few hundred days. The alternative, that a change in the radio emission associated with the ring impact is yet to be observed, may indicate an acceleration time greater than $\\sim 1800$ days, the interval from 1995 February to 2000 May. The model for the radio reappearance of \\SN\\ in 1990 July (BK92; DBK95) suggested that the acceleration time was $\\sim 300\\;\\rm days$.  Theory suggests that the time required for diffusive shock acceleration scales as $(B v_s^2)^{-1}$ where $B$ is the magnetic field and $v_s$ is the shock speed. Given that the shock speed is most probably decreasing with time, the first scenario suggests a significant increase in $B$ between the first circumstellar density jump and the ring, while the latter suggests a decrease. If the field is that of the stellar wind, an increase in $B$ with radial distance occurs if the wind is confined (as it is in \\SN), but more detailed modelling is required to address such issues adequately."
    },
    "0011/astro-ph0011080_arXiv.txt": {
        "abstract": "We present results of long-term daily monitoring of a sample of Galactic radio-emitting X-ray binaries showing relativistic jets (RJXRB): SS433, Cyg~X-3, and GRS~1915+105, with the RATAN-600 radio telescope in the 0.6--22~GHz range.  We carried out the modeling calculations to understand the temporal (1--100 days) and spectral (1--22~GHz) dependence. We tested the finite jet segment models and we found that the geometry of the conical hollow jets is responsible for either a power law or an exponential decay of the flares.  SS433 was monitored for 100 days in 1997 and 120 days in 1999. From the quiescent radio light curves, we obtained clear evidence of a 6.04-day 10--15\\% modulation.  Three powerful flares (up to 13~Jy) from Cyg X-3 were detected in April 2000. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011049_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recently, many X-ray pulsars have been discovered in the Small Magellanic Cloud (SMC), and now $\\sim 20$ X-ray pulsars are known in this galaxy. Most of them are interpreted as being X-ray binary pulsars (XBPs), because of their hard X-ray spectrum, long pulse period ($\\gtsim 1$~s), large flux variability, and/or the existence of a massive star counterpart (Yokogawa et al.\\ 2000 and references therein). Massive star counterparts have been mostly revealed to be a Be star by detailed optical observations on individual objects (Hughes, Smith 1994; Southwell, Charles 1996; Buckley et al.\\ 1997; Cowley et al.\\ 1997; Israel et al.\\ 1997; Coe et al.\\ 1998; Lamb et al.\\ 1999; Stevens et al.\\ 1999; Coe, Orosz 2000; Coe et al.\\ 2000), except for SMC X-1 with a B (not Be) supergiant star companion (Bildsten et al.\\ 1997).  Haberl and Sasaki (2000) investigated the spatial correlation between SMC X-ray sources and emission-line objects (hereafter ELOs) cataloged by Meyssonnier and Azzopardi (1993; MA93) as well as Murphy and Bessell (2000). They found that most of the already established Be/X-ray binaries have a counterpart in these ELO catalogs, and thus proposed that other X-ray sources having an ELO counterpart are also likely to be Be/X-ray binaries. They concluded that the number ratio of Be/X-ray binaries to OB supergiant X-ray binaries in the SMC is extremely higher than that in our Galaxy. They also pointed out that the main body of the SMC contains no OB supergiant X-ray binary, suggesting a different star-formation history between the main body and the eastern wing. In this context, it would be important to search for pulsations from X-ray sources with no ELO counterpart.   In this letter, we report on the discovery of coherent pulsations with ASCA from AX~J0057.4$-$7325 (Torii et al.\\ 2000), for which neither a Be star nor an ELO counterpart has been known. We also investigated several catalogs of optical stars to find a possible optical counterpart. ",
        "conclusions": "Six ROSAT observations and no Einstein observation have covered the position of AX~J0057.4$-$7325. To investigate the long-term flux variation of this source, we used the ROSAT archival data as follows. We derived the count rate from AX~J0057.4$-$7325 (= RX~J0057.3$-$7325) in each observation, after background subtraction. We assumed that $\\Gamma$ and $N_{\\rm H}$ were the same as those derived from the ASCA observations ($\\Gamma = 0.9$ and $N_{\\rm H} = 2\\times 10^{21}$ cm$^{-2}$), and converted the count rate to flux with the {\\tt PIMMS} software. We also made a vignetting correction according to the source's off-axis angle in each observation. The results are given in table 1: AX~J0057.4$-$7325 was detected in observations R1 and R2 with a signal-to-noise ratio of 2 and 10, respectively. By comparing the fluxes in observations R4 and A2, we conclude that the flux has changed by a factor of $\\gtsim 10$. The flux variability, the hard X-ray spectrum, and the long pulse period are all consistent with a scenario that AX~J0057.4$-$7325 is an XBP with a companion of either a Be, an OB supergiant, or a low-mass star.   As far as we have investigated, we found only one optical source, MACS~J0057$-$734\\#010, in the ASCA error circle. It is interesting that there exists no ELO, i.e., a Be star candidate, in the ASCA error circle. Because the nearest ELO, No.\\ 986 in MA93, is located $1.\\hspace{-2pt}'4$ away from AX~J0057.4$-$7325, a possibility for the optical counterpart would be ruled out. This is a rare case where an X-ray pulsar in the SMC is not associated with a Be star or an ELO.  {\\small \\begin{center} Table~1.\\hspace{4pt}Flux variability of AX~J0057.4$-$7325. \\\\ \\end{center} \\begin{tabular*}{0.5\\textwidth}{@{\\hspace{\\tabcolsep} \\extracolsep{\\fill}}p{2.5pc}ccc} \\hline\\hline\\\\ [-6pt] Obs.ID.&Date$^*$  &  Instrument  &  Flux$^\\dagger$         \\\\ &(MJD)     &              &  (erg~s$^{-1}$~cm$^{-2}$)      \\\\[4pt]\\hline\\\\[-6pt] R1\\dotfill& 48550.301 & ROSAT/PSPC   & $1.1\\times10^{-12}$ \\\\ R2\\dotfill& 48732.665 & ROSAT/PSPC   & $1.3\\times10^{-12}$ \\\\ R3\\dotfill& 49118.070 & ROSAT/PSPC   &$<3.5\\times10^{-13}$ \\\\ R4\\dotfill& 49298.538 & ROSAT/PSPC   &$<2.4\\times10^{-13}$ \\\\ R5\\dotfill& 49321.271 & ROSAT/PSPC   &$<3.9\\times10^{-13}$ \\\\ R6\\dotfill& 49470.832 & ROSAT/HRI    &$<1.1\\times10^{-12}$ \\\\ A1\\dotfill& 51310.141 & ASCA/GIS+SIS & $1.2\\times10^{-12}$ \\\\ A2\\dotfill& 51660.052 & ASCA/GIS     & $2.4\\times10^{-12}$ \\\\[4pt]\\hline \\end{tabular*} \\vspace{1pt}\\par\\noindent $*$ Middle of the observations. \\par\\noindent $\\dagger$ In the 0.7--10.0~keV band, after vignetting correction. }  We note that AX~J0057.4$-$7325 is located at the edge of the main body, fronting to the eastern wing. The fact that OB supergiant X-ray binaries in the SMC (only SMC X-1 and EXO 0114.6$-$7361) are both located in the eastern wing may lead us to suspect that AX~J0057.4$-$7325 would be the third example. Because the discovery of a true optical counterpart would provide essential information, we strongly encourage deep and detailed optical observations on this field. \\par \\vspace{1pc}\\par We thank Dr.\\ M.J.\\ Coe for his valuable suggestions about MACS~J0057$-$734\\#010 and optical counterparts of the SMC sources. We also thank Dr.\\ F.\\ Haberl for useful comments on the SMC PSPC catalog. The ROSAT data were obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. J.Y.\\ and T.K.\\ are supported by JSPS Research Fellowship for Young Scientists. We are grateful for constructive comments from an anonymous referee."
    },
    "0011/astro-ph0011339_arXiv.txt": {
        "abstract": "Using N-body numerical simulations, we examine possible mechanisms for the origin of the intracluster starlight features recently discovered in the Coma cluster of galaxies. We show that tidal interactions of a ``normal'' elliptical galaxy with a symmetrically distributed dark matter potential do not produce the observed intracluster starlight features. A head-on collision of two normal ellipticals does, however, explain the origin of the intracluster features with surface brightness about 26 mag/arcsec$^2$). Another possible explanation for the intracluster starlight features is galactic tidal interactions with massive intracluster substructures. The presence of substructures in Coma with a density of $\\sim$ 0.1 times the density in the center of the galaxy, and with a total mass about three orders of magnitude larger than the mass of a normal elliptical galaxy, would account for the observations. ",
        "introduction": "Gregg \\& West (1998) have recently reported the discovery of intracluster starlight features in the Coma cluster of galaxies. The features are revealed in R-band images of the Coma cluster, and are also seen in the B, V and I bands, confirming their existence. The plume-like features are tens of kiloparsecs in width, and extend to radii of more than a hundred kiloparsecs. The mean R-band surface brightness of the plumes is low, about 26 mag/arcsec$^2$, but their absolute magnitude $\\approx - 19.0$ mag is comparable to the absolute R-band magnitude of a normal elliptical galaxy. Possible explanations of the origin of the plumes invoke the idea of the disruption of galaxies by galaxy-galaxy interactions (Dubinski et al 1996, Moore et al 1996), or the tidal stripping of galaxies by the global cluster gravitational field (Merritt 1984). However, the stellar plumes observed in the Coma cluster can not be produced so easily. Using N-body simulations, we have studied the possibility of forming stellar plumes with surface brightnesses 6--7 magnitudes fainter than the surface brightness in the centers of the normal elliptical galaxies similar to those observed in the Coma cluster. We have found that the tidal interaction of galaxies with the gravitational field of the cluster core cannot be responsible for the origin of intracluster starlight features in Coma. A head-on collision of two normal elliptical galaxies with a relative velocity less than the cluster velocity dispersion, or tidal interactions with intracluster density inhomogeneities might explain the origin of stellar plumes with the observed properties. ",
        "conclusions": "We considered different possibilities for the formation of stellar plumes in clusters of galaxies, namely the tidal disruption of an elliptical galaxy by the cluster core, direct galactic collisions, and the interaction of an elliptical galaxy with hypothetical cluster substructures.  Collisions or tidal interactions often produce plumes of particles in numerical simulations. We need, however, a quantitative estimate of whether such plumes could be responsible for the observed stellar structures seen in Coma. To do this, we calculated the theoretical surface brightness in the plume assuming constant mass-to-luminosity ratios for the plume and the galaxy.  The characteristic surface brightness in the central regions of elliptical galaxies in the V-band is about 20 mag/arcsec$^2$ (Binney \\& Merrifield 1998), which for the R-band, gives a value of about 19.5 mag/arcsec$^2$ (Frei \\& Gunn 1994). We assumed that our model galaxy has the same surface brightness in central regions before an interaction, and that tidal interactions do not change the stellar content of the plume. These assumptions allow us to estimate surface brightness of resulting features in our numerical simulation by comparing corespondinc column densities of particles. The last assumption excludes from our consideration the possible role of star formation triggered in the plumes by tidal interactions.  Although tidally induced star formation in the plumes is likely to take place for the collisions of gas rich disk galaxies, this is not the case for the stellar plumes observed in Coma. According to Gregg \\& West (1998), the colors in the plumes are consistent with those of evolved stellar polulations observed in intermediate brightness normal elliptical galaxies.  The photometric uncertainties, and uncertainties of the removal of superimposed objects leave, however, open the possibility of significant star fromation in the plumes.  {\\it Direct galactic collisions}. Figure 1 shows the result of a direct collision of two equal-mass Plummer spheres taken at time $ 1.7\\times 10^8$ years after the beginning of the interaction. The initial separation of the galaxies and their initial relative velocity were taken to be 32 kpc and 408 km/sec. The collision forms a broad stellar plume connecting both galaxies, with the surface brightness in the plume about 26.5 mag/arcsec$^2$(Figure 2). The direct collisions of equal-mass galaxies might therefore constitute a possible explanation of the origin of the observed stellar plumes in Coma. The high velocity dispersion of about 1000 km/sec in Coma, however, makes this mechanism rather ineffective. Gregg \\& West (1998) argue that the galactic subgroups in Coma, with their relatively low internal velocity dispersions, will increase the effectiveness of the disruption of galaxies, and our simulations concur with this conclusion. The characteristic velocity dispersion within groups, which are fairly uniformly distributed in space and which show no preference for rich clusters, is about 250 km/sec (Hickson 1997). This value is much smaller than the typical velocity dispersions in rich clusters of galaxies. If a cluster is not virialized, then strong local interactions can occur within in-falling subgroups of galaxies (Conselice \\& Gallaher 1998). Burns et al (1994) present convincing arguments that such infalling groups of galaxies do indeed exist. They found that the galactic group NGC 4839 has fallen into the Coma Cluster, passing through the cluster core about 2 Gyr ago. The NGC 4839 galactic group cannot be responsible, however, for the observed starlight features in Coma. All four stellar features are located within $ \\sim 0.5$ Mpc of the cluster center, and cannot survive for more than 1-2 Gyr (Gregg \\& West 1998), which is shorter than the epoch of passage of the NGC 4839 group through the central regions of Coma.   {\\it Core tidal interactions}. The tidal interaction of galaxies with the gravitational potential of the cluster core is a likely explanation for the origin of stellar intracluster structures in Coma. We simulated an extreme case of such an interaction with a galaxy plunging into the cluster core. In these simulations, the cluster core radius adopted was 270 kpc, and the core density equal to $3\\times 10^{15} M_{\\odot}$ / Mpc$^3$ which are the extreme ``disruptive'' parameters of the mass distribution in Coma. The galaxy was placed 600 kpc away from the cluster center with zero initial relative velocity. Figure 3 shows surface density contour plots taken at time $T=3.36 \\times 10^9$ years after the beginning of the simulations, when the galaxy made one oscillation in the cluster core. The surface brightness in the patchy plume is 32-34 mag/arcsec$^2$, and apparently cannot explain the origin of the bright starlight features in Coma. This result is also illustrated by Figure 4, which shows the fraction of escaped particles as a function of time. A total of 10 \\% of the mass is stripped during the interaction time. This stripped mass cannot be regarded as the main source of intracluster stellar material in Coma.  {\\it Tidal disruption by the ``local core''}. One more possibility exists for the formation of the bright intracluster structures seen in Coma. The X-ray surface brightness measurements of Coma cluster reveal the regions of the enhanced X-ray emission with characteristic angular size less than the core radius of the cluster ( Soltan \\& Fabricant 1990). The intracluster stellar features in Coma may result from tidal interactions with small scale density fluctuations associated with the enhanced X-ray emission. Vikhlinin et al (1994) found indeed such mass concentrations in the core region of the Coma cluster. Both of the discovered density enhancements are centered at the positions of the brightest elliptical galaxies in the cluster: NGC 4874 and NGC 4889. Assuming that X-ray surface brightness enhancements result from the density enhancements of the underlying dark matter, Vikhlinin et al (1994) estimated the masses of inhomogeneities within 80 kpc to be about $10^{13} M_{\\odot}$. Remarkably, the most prominent plume-like structure, the LSB object 1 discovered by Gregg and West, is adjacent to NGC 4874. The core radius of the density enhancement centered at the position of NGC 4874 is about 100 kpc (Vikhlinin et al 1994) which gives the density of the ``local core''about $5 \\times 10^{-25}$ g/cm$^3$. This value is too low to disrupt noticeably a galaxy by tidal interaction.  The parameters of the density enhancement centered at the position of NGC 4889 are more favorable for the galactic disruption.  This density enhancement has core radius about 24 kpc, and the central density about $8 \\times 10^{-24}$ g/cm$^3$. We simulated an interaction of a galaxy with the NGC 4889 density enhancement. A galaxy was placed 200 kpc away from the density enhancement with the initial velocity 500 km/s. The galaxy was allowed then to collide non-centrally with the density enhancement with the impact parameter equal to 24 kpc. Figure 5 shows the result of such high-speed tidal interaction. Interaction with NGC 4889 density inhomogeneity forms a plume with the surface brightness only 29-30 mag/arcsec$^2$. This value is too low to account for the properties of the observed plumes. The density enhancements, however increase an effective cross-section of the direct collision with  the giant cD galaxies nested in their central regions. This is a likely scenario of the formation of the most luminous plume LSB no 1 adjusted to NGC 4884.  Stars in the plume structures which are induced in direct galactic collisions have higher velocity dispersion compared to the plumes produced in tidal interactions. The collisional plumes live therefore shorter than tidally induced structures. Our numerical simulations show that the density inhomogeneities tidally disrupt an elliptical galaxy if the central density of a ``local core'' is about 10 percent of the central galactic density. Such value can be achieved if an inhomogeneity has total mass about one thousand of a mass of a normal elliptical galaxy, and a core radius $\\le$ 15 kpc. If such inhomogeneities of the intracluster matter do really exist, then the starlight features in clusters of galaxies give an opportunity to probe the cluster density fluctuations."
    },
    "0011/astro-ph0011425_arXiv.txt": {
        "abstract": "Thermal conduction in tangled magnetic fields is reduced because heat conducting electrons must travel along the field lines longer distances between hot and cold regions of space than if there were no fields. We consider the case when the tangled magnetic field has a weak homogeneous component. We examine two simple models for temperature in clusters of galaxies: a time-independent model and a time-dependent one. We find that the actual value of the effective thermal conductivity in tangled magnetic fields depends on how it is defined for a particular astrophysical problem. Our final conclusion is that the heat conduction never totally suppressed but is usually important in the central regions of galaxy clusters, and therefore, it should not be neglected. ",
        "introduction": "In order to solve the problem of electron thermal conduction in a stochastic magnetic field, one should consider separately two effects that reduce the conduction (Pistinner \\& Shaviv 1996; Chandran \\& Cowley 1998). The first effect is that the heat conducting electrons have to travel along tangled magnetic field lines, and therefore, they have to travel longer distances between hot and cold regions of space (Tribble 1989, Tao 1995). The second effect is that electrons, while they are traveling along the field lines, become trapped and detrapped between magnetic mirrors, regions of strong magnetic field (Chandran, Cowley, \\& Ivanushkina 1999).  In our recent paper (Malyshkin \\& Kulsrud 2000) we studied the second effect, and we found the reduction of the thermal conductivity parallel to the magnetic field lines, $\\kappa_\\parallel$, relative to the Spitzer value for the thermal conductivity, $\\kappa_{\\rm S}$, caused by the presence of magnetic mirrors. In this paper we consider the first effect, and we calculate the further reduction of thermal conduction relative to $\\kappa_\\parallel$, caused by the tangled structure of the magnetic field lines. As a result, we obtain the total reduction caused by the both effects, and we calculate the effective thermal conductivity $\\kappa_{\\rm eff}$, which one has to apply for astrophysical problems. It turns out that there is no an universal definition of $\\kappa_{\\rm eff}$, and the result depends on the particular astrophysical problem under consideration.  The structure of this paper is the following. In the next section we consider the random walk of tangled magnetic field lines by making use of a diffusion approximation for this random walk. We find the expressions for the probability distributions of the field line lengths and for the Laplace transform, in the field line length, of these distributions. In Section~\\ref{MODEL} we use the Laplace transforms to calculate the effective thermal conductivity for our stationary model of a galaxy cluster. In Section~\\ref{T_EVOLUTION} we consider a time-dependent model and study the time evolution of temperature in a galaxy cluster. Finally, in Section~\\ref{DISCUSSION} we discuss our results, compare them with the results reported in previous papers, and consider the applications to cooling flow problem. For two opposite limiting cases our results for $\\kappa_{\\rm eff}$ are: one is similar to that of Tribble (1989), and the other coincides to that of Tao (1995). ",
        "conclusions": "\\label{DISCUSSION}  We would like to start the discussion with by stressing one of the main results we found: the actual value of this effective conductivity in the tangled magnetic fields depends on how this effective conductivity is defined. In our paper we used two ``natural'' definitions of $\\kappa_{\\rm eff}$ for two very simple models of a cluster of galaxies: a time-independent model in Section~\\ref{MODEL} and a time-dependent model in Section~\\ref{T_EVOLUTION}. The results for $\\kappa_{\\rm eff}$, reported in figures~\\ref{FIG_KAPPA_EFF}(a) and~\\ref{FIG_KAPPA_EFF}(b), are the same only provided $\\beta\\gg 1$, and they are significantly different (by up to a factor of ten) when $\\beta\\simlt 1$. We conclude that there does not exist an universal result for the effective conductivity in the tangled magnetic fields. One has to define $\\kappa_{\\rm eff}$ and to calculate it for a particular astrophysical problem that he/she considers. We believe that such calculation is possible for many problems, including numerical simulations of galaxy cluster formation and of cooling flows, by making use of our diffusion approximation method for the random walk of tangled magnetic field lines.  It is useful to compare our results for the effective conductivity in the tangled magnetic fields with those reported in previous papers. As we said in the last paragraph of Section~\\ref{MODEL}, in the limit $\\beta\\gg 1$ our result for $\\kappa_{\\rm eff}$ coincides with the result obtained by Tao (1995). This simple formula (for $\\beta$ large) corresponds to lines being in order $X_0/\\epsilon$ in length, a result sometimes believed (remember, that $\\epsilon\\ll 1$ is the ration of the field mean component to the field random component). But in fact, it is only valid as a lower limit on $\\kappa_{\\rm eff}$, except for large $\\beta$. However, we can speculate that it revolves around the very definition of $\\epsilon$. We give $\\epsilon$ first and then impose the statistics of the field. If the statistics is given first and the mean field is defined afterwards, as the rms value of the mean field on the scales $\\sim X_0$, then $\\kappa_{\\rm eff}$ would be closer to our results.  In the limit $\\beta\\ll 1$, our result, $\\kappa_{\\rm eff}\\sim 0.1 \\kappa_{\\parallel} l_0^2/X_0^2$, is consistent with the result of Tribble (1989).~\\footnote{ Tribble has $\\kappa_{\\rm eff}=4\\kappa_{\\parallel} l_0^2/X_0^2$, see his equation~(11). The difference arises because: first, our boundary conditions~(\\ref{BOUNDARY_MINUS}) and~(\\ref{BOUNDARY_PLUS}) for the random walk of field lines are different from those of Tribble; second, the magnetic field lines in our model are allowed to random walk in three dimensions, while Tribble considered a restricted one-dimensional random walk. } To the best of our knowledge, there have no results for $\\kappa_{\\rm eff}$ when $\\beta\\sim 1$ obtained before.  Recently Chandran and Cowley (1998) have suggested that the diffusion of the heat conducting electrons perpendicular to the magnetic field lines is crucially important for the heat transport. In their model the statistically independent random step by the electrons is equal to $\\sqrt{2L_{\\rm RR}l_0}$, where $L_{\\rm RR}\\sim l_0\\ln{(l_0/\\rho_e)}$ is the Rechester-Rosenbluth length. This length is the distance along magnetic field lines over which a separation between two initially neighboring lines grows exponentially from the electron gyro-radius $\\rho_e$ up to the field decorrelation length $l_0$ (because of the Kolmogorov-Lyapunov exponential divergence of field lines). Chandran and Cowley suggested that the perpendicular diffusion of electrons and this exponential divergence of the magnetic field lines enable the electrons to conduct heat between different field lines, which may enhance the effective thermal conductivity considerably. However, we believe that the picture considered by Chandran and Cowley is rigorously valid only if electrons can move a distance $L_{\\rm RR}$ along field lines without collisions. In fact, there are collisions, and if the electron mean free path $\\lambda$ is less than the Rechester-Rosenbluth length $L_{\\rm RR}$, the exponential divergence of field lines does not help electrons to diffuse perpendicular field lines. In clusters of galaxies, the ratio of $L_{\\rm RR}$ to $\\lambda$ is $L_{\\rm RR}/\\lambda \\approx 5000\\,(l_0/10\\,{\\rm kpc})\\,(T/10^7{\\rm K})^{-2}\\,(n/10^{-3}{\\rm cm}^{-3})$. As a result, the perpendicular heat transport can generally be neglected (we accept $3$--$10\\,{\\rm kpc}$ as a typical scale for magnetic fields in clusters of galaxies, see Kronberg 1994, Eilek 1999). The Chandran and Cowley's model can be very important for some other astrophysical problems, but the consideration of their model is beyond of the scope of this paper.  Finally, let estimate the importance of the heat conduction for the formation and evolution of galaxy clusters, using figures~\\ref{FIG_T_MAP}. The parameter $\\beta$ and the dimensional time $\\tau$, given by equations~(\\ref{BETA}) and~(\\ref{TAU}), can be estimated as \\beq \\beta=20\\,\\frac{X_0}{1\\,{\\rm Mpc}}\\,\\frac{10\\,{\\rm kpc}}{l_0}\\,\\frac{\\epsilon\\cos\\theta}{0.1}, \\quad \\tau=5\\!\\times\\!10^{-6}\\,\\frac{\\kappa_\\parallel}{\\kappa_{\\rm S}}\\,\\frac{t}{10^{10}{\\rm yrs}}\\, \\frac{10^{-3}{\\rm cm}^3}{n}\\left[\\frac{T}{10^7K}\\right]^{5/2} \\left[\\frac{l_0}{10\\,{\\rm kpc}}\\right]^2\\left[\\frac{1\\,{\\rm Mpc}}{X_0}\\right]^4. \\nonumber \\eeq The reduction of the parallel thermal conductivity relative to the Spitzer value, caused by magnetic mirrors, depends on the ratio $l_0/\\lambda=160\\,(l_0/10\\,{\\rm kpc})\\,(T/10^7{\\rm K})^{-2}\\,(n/10^{-3}{\\rm cm}^{-3})$ [Malyshkin \\& Kulsrud 2000]. The typical values are $\\kappa_\\parallel/\\kappa_{\\rm S}\\sim 1/10$ for the hot and low-density halo of a galaxy cluster, and $\\kappa_\\parallel/\\kappa_{\\rm S}\\sim 1$ for the cluster central region. We see that for the halo, where $X_0\\sim 1\\,{\\rm Mpc}$, $l_0\\sim 10\\,{\\rm kpc}$, $T\\sim 5\\times 10^7\\,{\\rm K}$ and $n\\sim 2\\times 10^{-4}{\\rm cm}^{-3}$, we have $\\beta\\sim 20$ and $\\tau\\sim 10^{-4}$. Thus, according to figures~\\ref{FIG_T_MAP}, the heat conduction is unimportant there, while according to figures~\\ref{FIG_KAPPA_EFF}, the effective thermal conductivity is reduced by a factor of $\\sim 2000$. At the same time, for the cluster central region we have $X_0\\sim 0.1\\,{\\rm Mpc}$, $l_0\\sim 3\\,{\\rm kpc}$, $T\\sim 2\\times 10^7\\,{\\rm K}$, $n\\sim 10^{-3}\\,{\\rm cm}^{-3}$, $\\beta\\sim 7$ and $\\tau\\sim 0.03$, so the heat conduction is important there, and the effective thermal conductivity is reduced by a factor of $\\sim 200$. This conclusion agrees with those of Tau (1995) and of Rosner and Tucker (1989). Thus, with certain parameters, such as these, thermal conduction can be important even with such a large reduction factor. Figures~\\ref{FIG_T_MAP}(a) and~\\ref{FIG_T_MAP}(b) give the direct temperature evolution without reference to reduction factors, and perhaps because of the sensitivity of the reduction factor to its definition, they are probably the more useful expressions of thermal evolution.  Suginohara and Ostriker in their hydrodynamic simulations of galaxy cluster formation encountered a  ``cooling catastrophe'', which appears as a steep non-realistic rise in the density profile of the relaxed core of a galaxy cluster because of the fast cooling in the core (Suginohara \\& Ostriker 1998). The heat conduction into the core was among solutions suggested by them. However, they believed that the conduction can be neglected if it is reduced by a factor of $30$ or more. In this paper we find that the thermal conduction is very important in the cluster central region, despite the fact that in our models it is reduced by a factor of $\\sim 200$. Thus, the conduction should be included in hydrodynamic simulations of galaxy clusters, even when the reduction factor is large."
    },
    "0011/astro-ph0011154_arXiv.txt": {
        "abstract": "We apply the synchrotron proton blazar (SPB) model to the April 1997 flare of Markarian 501 and find we are able to fit the observed spectral energy distribution.  We explore the effect of target photon density on the high energy part of the spectral energy distribution (SED) for fixed assumed magnetic field, emission region size and Doppler factor and find that the luminosity and peak frequency of the high energy part of the SED may depend on the luminosity of the low energy part of the SED in high-frequency peaked BL Lac objects (HBL). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011362_arXiv.txt": {
        "abstract": "We present the results of a spectral line survey from 138.3 to 150.7 GHz toward Orion-KL. The observations were made using the 14 m radio telescope of Taeduk Radio Astronomy Observatory. Typical system temperatures were between 500 and 700 K, with the sensitivity between $0.02 - 0.06$ K in units of $\\rm T_A^*$.  A total of 149 line spectra are detected in this survey. Fifty lines have been previously reported, however we find 99  new detections. Among these new lines, 32 are  `unidentified', while 67 are from molecular transitions with known identifications. There is no detection of H or He recombination lines. The identified spectra are from a total of 16 molecular species and their isotopic variants. In the range from 138.3 to 150.7 GHz, the strongest spectral line is the J=$3-2$ transition of CS molecule, followed by transitions of the $\\rm H_2CO$, $\\rm CH_3OH$, $\\rm CH_3CN$, and $\\rm SO_2$. Spectral lines from the large organic molecules such as $\\rm CH_3OH$, $\\rm CH_3OCH_3$, $\\rm HCOOCH_3$, $\\rm C_2H_5CN$ and $\\rm CH_3CN$ are prominent; with 80 \\% of the identified lines arising from transitions of these molecules. The rotational temperatures and column densities are derived using the standard rotation diagram analysis for $\\rm CH_3OH$ ($\\rm ^{13}CH_3OH$), $\\rm HCOOCH_3$, $\\rm CH_3CN$ and $\\rm SO_2$ with $\\rm 10\\sim 270~K$ and $\\rm 0.2\\sim 20\\times 10^{15}~cm^{-2}$. These estimates are fairly comparable to the values for the same molecule in other frequency regions by other studies. ",
        "introduction": "Orion-KL is a unique source for study of numerous astro-chemical aspects associated with massive star formation. Spectral line surveys can provide an unbiased view on molecular constituents of such regions by fully probing various spectral lines in wide frequency bands, and also are necessary to understand physical and chemical processes in a massive star-forming region.  Toward Orion-KL there have been numerous molecular line surveys in the spectral ranges of $70 - 360$ GHz and $607 - 900$ GHz by other researchers (see Table 1). However, there are no spectral line surveys in the 2 mm band except for $150 - 160$ GHz by Ziurys \\& McGonagle (1993, hereafter ZM93). ZM93 have detected over 180 spectral lines including about 45 unidentified features, and so concluded that Orion-KL is also very rich region for molecular lines in 2 mm band. We have been carrying out a spectral line survey toward Orion-KL to cover the entire `$\\it unexplored$' 2 mm spectral band. This paper presents our observational results of the survey from 138.3 to 150.7 GHz using the 14 m telescope of Taeduk Radio Astronomy Observatory (TRAO). ",
        "conclusions": "A total of 149 spectral lines are detected with a complete spectrum shown in Fig. 1. Detailed identification of the molecular emission lines, LTE rotation diagram analysis of some molecular species, other interesting lines, and spatial origin of the molecular species are described in this section.  \\subsection{Identification of detected lines} The rest frequency scale for identification of the detected molecular emission lines is obtained by  assuming $\\rm V_{LSR} = 8.4~km~s^{-1}$. This value is determined from a Gaussian fit to the CS J$=2-1$ spectrum. This can be somewhat arbitrary as the reference value, because Orion-KL is known to consist of a number of kinematically different components with sizes of a few $10''$, such as extended ridge, hot core, compact ridge, and plateau (Sutton et al. 1995). The existence of these different velocity components may yield an uncertainty in the derived frequency of $\\sim$ 1 MHz or slightly larger.  We used two catalogs for line identification: the Lovas catalog (Lovas 1992) which is a collection of the molecular line frequencies determined from astronomical observations, and the JPL catalog (Pickett et al. 1998), a collection of theoretical calculations of molecular lines. In addition, catalogs by Anderson et al. (1992), Xu \\& Lovas (1997), and Tsunekawa et al. (2000, in private communication), and the Kawaguchi catalog (Kawaguchi 1999, in private communication) are complementarily used  for the line identification.  If a detected line matches the frequency of any identified line in Lovas catalog to within $\\sim \\pm$ 1 MHz, it is regarded as previously `observed'. However, if a detected line is not in the Lovas catalog, but in the list of one of the other five catalogs, it is considered as a `new' detection. This case requires that any other transitions of the species responsible for the `new' line  should have been previously detected from other astronomical observations. We mark `NDT' in the 6th column of Table 2 for the `newly' detected transition.  If a detected line is not listed in any catalog, we designate it as an unidentified or ``U'' line.  Occasionally a high excitation transition of a known species matches the frequency of a detected line, although no lower excitation transitions of the species have been detected. These matches are regarded as coincidences and marked as unidentified.  With a detection limit of $\\sim 3\\sigma$, a total of 149 molecular line spectra are detected. The detailed identifications of the spectra are given in Table 2. The frequency in the first column is the value given from the catalogs provided in the 5th column. For the `U' lines, we give the frequencies determined from our observations. Molecular species, their transitions, intensity in $\\rm T_A^*$, and special comments for the spectra are given in the 2nd, 3rd, 4th, and 6th column of the Table 2.  Fifty of 149 spectral lines are found to  have been previously identified from other studies, while there are 99  lines which represent the first detections towards an astronomical source. Sixty-seven of the new 99  lines are transitions predicted from theoretical calculations or detected from laboratory measurements. However, 32 detected lines remain ``unidentified''. We display all scans with identifications of the spectra in Fig. 2.  The spectral region between $150000$ MHz $\\sim 150760$ MHz overlapped with the survey region by ZM93. ZM93 detect $\\sim$19 spectral lines in this region while we detect only 12 lines.  The sensitivity  between the two surveys is similar and the antenna temperatures for the strong lines (e.g., transitions of $\\rm CH_3OH$, $\\rm SO_2$, and $\\rm H_2CO$) for both observations are comparable. Moreover, the 2mm beam efficiency of the two telescopes is about the same ($\\sim0.4$). The identifications for strong eight lines  are the same between two studies, but for weak lines, different identifications are given. We especially note two transitions of $\\rm CH_3OCH_3$ at 150163.1 MHz and 150467.3 MHz. The former 150163.1 MHz line is considered to be `unidentified' by ZM93 while the latter 150467.3 MHz line is not on their list. ZM93 assign the spectra at 150415.1 MHz and 150449.2 MHz as transitions of $\\rm CH_3OH$ and $\\rm HCOOCH_3$, respectively. However, the line intensities in both observations are close to the noise level and so we leave these as undetected.  We also find two U lines in this overlapped spectral region for which ZM93 did not give any identification. In contrast ZM93 find 4 U lines for  which we did not provide an identification. However, the antenna temperatures of the six U lines are $\\sim 0.1$ K, which is close to noise level. Follow-up observations would be necessary for confirmation of the identification of the weak lines.  The identified 50 emission lines are found to originate from a total of 16 species and their isotopic variants. The strongest molecular transition between 138.3 to 150.7 GHz is the J=$3-2$ transition of the CS molecule, followed by the lines of the $\\rm H_2CO$, $\\rm CH_3OH$, $\\rm CH_3CN$, and $\\rm SO_2$.  An interesting point is that a considerable number of the detected lines are from large organic molecules such as $\\rm CH_3OH$ (30), $\\rm ^{13}CH_3OH$ (4), $\\rm CH_3OCH_3$ (20), $\\rm HCOOCH_3$ (16), $\\rm C_2H_5CN$ (15), and $\\rm CH_3CN$ (7) (the numeric in parenthesis is the number of the detected transitions of the molecule). About 80 \\% of the identified lines are from the transitions of these organic molecules.  \\subsection{The LTE Rotation Diagram Analysis and Individual Molecules} We perform the standard LTE rotational diagram analysis to derive rotation temperatures ($\\rm T_{rot}$) and column densities ($\\rm N$) of several molecules for which more than four transitions are detected. The derived physical quantities are compared with results by other studies in different wavelength bands toward Orion-KL.  The equation for the LTE rotation diagram analysis (e.g., Turner 1991) is given as $$ log L = log {{3k \\int T_A^* dv}\\over{8 \\eta_B \\pi^3 \\nu S\\mu_i^2 g_I g_k}} = log{N\\over Q_{rot}} -{{E_u~log~e} \\over {k~T_{rot}}}, \\eqno (1) $$ where k is the Boltzmann constant, $\\int T_A^* dv$ the integrated intensity, $\\eta_B$ the beam efficiency of the telescope, $\\nu$ the rest frequency of the spectrum, S the line strength, $\\mu_i$ the relevant dipole moment, $g_I$ the reduced nuclear spin weight, $g_k$ the K-level degeneracy, $E_u$ the upper state energy of the transition, $Q_{rot}$ the rotational partition function, and N the column density.  The values for  $g_I$, $g_k$, $E_u/k$, $S\\mu_i^2$, $\\int T_A^* dv$, and $log L$ for each molecule and their references are given in Table $3 - 10$. The partition functions $Q_{rot}$ are ${1 \\over 2} [{{\\pi (kT_{rot})^3} \\over {h^3ABC}}]^{1/2}$ for $\\rm H_2CO$ and $\\rm SO_2$, ${1 \\over 3} [{{\\pi (kT_{rot})^3} \\over {h^3AB^2}}]^{1/2}$ for $\\rm CH_3CN$, $ [{{\\pi (kT_{rot})^3} \\over {h^3ABC}}]^{1/2}$ for $\\rm CH_3OCH_3$ and $\\rm C_2H_5CN$, and $2 [{{\\pi (kT_{rot})^3} \\over {h^3ABC}}]^{1/2}$ for $\\rm CH_3OH$ and $\\rm HCOOCH_3$, where A, B, and  C are the rotational constants (Blake et al. 1986; Turner 1991).  The above equation is derived under the assumption that all lines are optically thin, all level populations are characterized by a single $T_{rot}$ (LTE), $\\rm T_{rot} >> T_{bg}$ (background temperature), and that the Rayleigh-Jeans approximation is valid for all transitions (see Turner 1991). If these assumptions are valid and the line identifications are correct, the distribution of $E_u/k$ versus log L of the data would exhibit a linear correlation, and the rotational temperature and the column density can be determined from the slope [$-(log e)/T_{rot}$] and intercept [$log (N/Q_{rot})$ at $E_u = 0$] through a linear least squares fit of the data.  On the other hand, if the distribution of $E_u/k$ versus log L of the data in the rotation diagram deviates greatly from the linear correlation, the assumptions of low optical depth and LTE may not be valid and requires further consideration in deriving $\\rm T_{ex}$ and N from the data. Moreover an incorrect line identification may also result in a large  deviation of the data point from the correlation. Therefore, the LTE rotational diagram analysis itself can be useful for checking secure identification of spectral lines.  We applied this analysis for 7 molecules and one isotopic variant: $\\rm CH_3OH$ ($\\rm ^{13}CH_3OH$), $\\rm HCOOCH_3$, $\\rm CH_3OCH_3$, $\\rm C_2H_5CN$, $\\rm SO_2$, $\\rm CH_3CN$, and $\\rm H_2CO$. The transitions of 4 molecules [$\\rm CH_3OH$, $\\rm HCOOCH_3$, $\\rm SO_2$, and $\\rm CH_3CN$] and one isotopic variant ($\\rm ^{13}CH_3OH$) are found to exhibit a linear distribution in the rotation diagram, allowing us to introduce the estimates of the physical quantities of the molecules and compare our results with those from other surveys. On the other hand,  the rotational diagrams for three molecules: $\\rm CH_3OCH_3$, $\\rm C_2H_5CN$ and $\\rm H_2CO$ show scattered non-linear distributions. Therefore, the rotation temperatures and the column densities for these molecules are not derived.  Fig. 3 shows the diagrams for 6 molecules and one isotope variant used in the analysis. The diagram for $\\rm H_2CO$  is not shown because of the large scatter seen in the data. In the following paragraphs we discuss the results of the rotational diagram analysis in sequence of increasing number of detected transitions of the molecule.  {\\bf CH$_3$OH--} Methanol is one of most widely observed species in molecular clouds and star-forming regions. The number of transitions detected (30) is the largest in our observations (Table 3). We find 19 new detections from a total of 30 methanol lines identified in our survey. The data distribution in the rotation diagram shows a fairly linear distribution with some scatter. In the rotation diagram analysis we use 26 transitions for which $\\rm E_u/k$ and $\\rm S\\mu^2$ are known. We removed a single data point from the [$3(1)-2(1)$ A-] transition in the analysis which shows a large deviation in the diagram (Fig. 3). The linear least squares fit to the data indicates a fairly good correlation (correlation coefficient; $r=-0.86$), yielding  results of $\\rm T_{ex}=160_{-17}^{+21}$ K and $\\rm N=2.1_{-0.7}^{+1.1}\\times 10^{16}~cm^{-2}$ which are close to 3 mm results ($\\rm T_{ex}=192$ K and $\\rm N=1.6\\times 10^{16}~cm^{-2}$) obtained by Turner (1991).  Four transitions of the rare isotopic species, $\\rm ^{13}CH_3OH$, are also detected with three new detections (Table 4). The linear least squares fit to the data, excluding a single deviated point [3(-1)-2(-1) E] results in a nearly perfect correlation ($r=-0.998$) (Fig. 3), producing $\\rm T_{ex}=40_{-2}^{+3}$ K and $\\rm N=1.2_{-0.2}^{+0.2}\\times 10^{14}~cm^{-2}$. Using the terrestrial value of $\\rm ^{12}C/^{13}C=89$ may provide an alternate (and more reliable) estimate ($\\rm N=1.1\\times 10^{16}~cm^{-2}$) of the column density of Methanol, which is close to the value determined using the LTE rotational method within the uncertainty.  {\\bf CH$_3$OCH$_3$--} Dimethyl ether is one of a few interstellar molecules with two internal rotors and we detect twenty transitions from this molecule in our survey (Table 5). In contrast to Methanol, there are no newly identified lines. The rotation diagram shows very scattered distribution of the data, indicating no  linear correlation between $\\rm E_u/k$ and log L (Fig. 3). This prevents us from inferring confident values of $\\rm T_{rot}$ and N. Note that Johansson et al. (1984), Blake et al (1986), and Turner (1991) have deduced $\\rm T_{ex}=63 - 91$ K and $\\rm N=1.3 - 3.0 \\times 10^{15}~cm^{-2}$ from their 3 mm data.  {\\bf HCOOCH$_3$--} Sixteen transitions of Methyl formate are detected (Table 6), with 8 lines believed to be `first' detections. Turner (1991) has shown that this species has two components in the rotation diagram and our data in the rotational diagram also seem to exhibit two components (Fig. 3). The rotation analysis for one component yields $\\rm T_{ex}=21.8_{-4.4}^{+7.2}$ K and $\\rm N=2.3_{-1.5}^{+4.9} \\times 10^{15}~cm^{-2}$. This rotational temperature is very close to that found by Turner (1991), but the column density is one order of magnitude higher than his result. The other, cooler component, shows $\\rm T_{ex}=18.1_{-3.6}^{+6.0}$ K and $\\rm N=2.6_{-1.7}^{+5.5} \\times 10^{14}~cm^{-2}$ has not been seen by any other studies.  {\\bf C$_2$H$_5$CN--} Fifteen transitions of Ethyl cyanide are detected (Table 7). Twelve lines are considered to be new detections. Interestingly, seven transitions seem to be K-type doubler lines whose frequencies are the same. The distribution of the data is not a simple linear appearance. So, the rotational diagram method does not seem to be applicable to derive $\\rm T_{ex}$ and N of this molecule.  {\\bf SO$_2$--} Sulfur dioxide is detected in seven transitions (Table 8) of which four are newly detected lines. Five strong lines are found to consist of two clear velocity components, a relatively narrow and bright component of $\\rm V_{LSR}\\approx 6 - 8~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 14 - 31~km~s^{-1}$, and a wide and weak component of $\\rm V_{LSR}\\approx 10 - 18~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 15 - 34~km~s^{-1}$. The data  of  Log L and $\\rm E_u/k$ for the former component show a good correlation ($r=-0.92$) with a slight scatter in the rotation diagram (Fig. 3). The analysis yields $\\rm T_{ex}=120_{-20}^{+21}$ K and $\\rm N=1.7_{-0.1}^{+1.5} \\times 10^{16}~cm^{-2}$. Considering the uncertainty, our $\\rm T_{ex}$ is fairly comparable to the value (138.5 K) found in the 3 mm observations by Turner (1991) and the value ($\\sim 140$ K) in 2 mm observations obtained by ZM93. The column density is also in close agreement with the results ($\\rm N\\approx 1.5\\times10^{16}~cm^{-2}$) by Turner (1991) and ($\\rm N\\approx 2.0\\times10^{16}~cm^{-2}$) by ZM93. The rotation diagram for the wide and weak component which we did not include in the Fig. 3 produces a very scattered distribution and we do not to derive $\\rm T_{ex}$ and the N.  {\\bf CH$_3$CN--} Methyl cyanide emissions  are detected in the seven K-components from the J=$8-7$ transition over very narrow frequency range of about 700 MHz (Table 9). One transition [8(6)-7(6)] is a new detection. Previously different investigators have suggested several different values for the  $\\rm T_{ex}$ and N primarily due to the large scatter in the data. The overall distribution of our data points in the rotation diagram shows some scatter in the correlation ($r=-0.61$) (Fig. 3). The least squares fit results in  $\\rm T_{ex}=273_{-101}^{+381}$ K and $\\rm N=1.3_{-0.1}^{+5.1} \\times 10^{15}~cm^{-2}$. The $\\rm T_{ex}$ is virtually the same as the value (274 K) for the hot core component by Blake et al. (1986) and the column density for the component is close to their estimate.  {\\bf H$_2$CO--} Formaldehyde has three detected transitions (Table 10). All of the transitions have been previously observed by other studies. The line shapes of all spectra clearly exhibit two components, one a narrow bright emission component   of $\\rm V_{LSR}\\approx 8.6 - 8.9~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 5 - 6 ~km~s^{-1}$, and the other a broad faint emission component (presumably plateau emission) of $\\rm V_{LSR}\\approx 7.0 - 8.3~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 15 - 20 ~km~s^{-1}$. Due to the high degree of scatter in the correlation we have not included the rotation diagram in Fig. 3 and also do not derive a rotational temperature and total column density for the two components.  In summary, the LTE rotational diagram analysis is fairly useful to derive important physical quantities ($\\rm T_{ex}$ and N) from four molecules and one isotopic variant, each with more than 4 detected transitions. All the values are comparable to the estimates in different frequency regions by other researchers. The estimated values are compared  with those from other studies in Table 11.  \\subsection{Other Interesting Lines}  {\\bf CS Lines--} Of the observed emission lines the CS molecule  stands out in that the CS (3$-$2) emission is the strongest among the detected transitions and its three rare isotopes such as $\\rm ^{13}CS$, $\\rm  C^{34}S$, and $\\rm  C^{33}S $ are also strongly detected in our survey (Table 12). The main isotope CS ($3-2$) line appears to contain several components, among which a bright ridge component of $\\rm V_{LSR}\\approx 5.5~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 8.2~km~s^{-1}$, and a broad plateau component of $\\rm V_{LSR}\\approx 6.5~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 21.5~km~s^{-1}$ are included. The three rare-isotope species seem to trace a ridge component of $\\rm V_{LSR}\\approx 8\\sim 9~km~s^{-1}$ and $\\rm \\Delta V_{FWHM}\\approx 5~km~s^{-1}$. Comparison of the CS emission with emissions from the isotope variants can give the CS optical depth. The ratios of integrated intensities between $\\rm  C^{32}S$ and $\\rm  C^{33}S$, and between $\\rm  C^{32}S$ and $\\rm  C^{34}S$ are about 36.6 and 5.5, respectively. These are significantly different from the terrestrial values of $\\rm ^{32}S/^{33}S=80$ and $\\rm ^{32}S/^{34}S=18$, meaning that the CS (3$-$2) emission is optically thick. This is also inferred from the fact that the ratio (about 11.9) of the integrated intensities between CS and $\\rm ^{13}CS$  is small compared with the terrestrial value of $\\rm ^{12}C/^{13}C=89$. Assuming $\\rm T_{ex}=65$ K and $\\rm ^{32}S/^{34}S=18$ which are adopted from Sutton et al. (1995), we obtain $\\rm N(C^{34}S)\\approx 1.1\\times 10^{13}~cm^{-2}$ and $\\rm N(CS)\\approx 2.0\\times 10^{14} ~cm^{-2}$ using the LTE method (e.g., Lee, Minh, \\& Irvine 1993). These estimates are one order of magnitude smaller than those of Sutton et al. (1995). This difference may be the result of the larger beam dilution in our survey.  {\\bf Recombination Lines--} The recombination lines in our survey frequency region are $\\rm H35\\alpha$ (147046.9 MHz), $\\rm H44\\beta$ (144474.0 MHz), $\\rm H50\\gamma$ (144678.9 MHz), $\\rm He35\\alpha$ (147106.8 MHz), $\\rm He44\\beta$ (144532.9 MHz), and $\\rm He50\\gamma$ (144737.9 MHz). The most plausible candidate of the radio recombination lines may be the U147112.9 line which has close frequency to $\\rm He35\\alpha$, but the frequency difference between two is somewhat large ($\\sim 6$ MHz). So we leave the U147112.9 line unidentified. There is no clear evidence of detection for other recombination lines in our survey.  {\\bf Unidentified (U) Lines--} A total of 32 U lines toward Orion-KL are found in our survey. Several lines (e.g., U143006.7 and U143263.3) are tentative detections because of low S/N ratios of the spectra. We note that U139561.9, U146372.4, U146622.4, U146984.5, U147759.8, and U147943.7 are quite bright ($\\rm T_A^* \\ga 0.3~K$). The frequency, $\\rm \\Delta V_{FWHM}$, antenna temperature, and integrated intensity of the U lines are given in Table 13 and marked with `U' on the spectra in Fig. 2. All U lines have $\\rm \\Delta V_{FWHM}\\approx 2 - 18~km~s^{-1}$ and seem to originate from various components such as ridge, hot core, and plateau. The U147943.7 line has two velocity components, presumably a hot core component ($\\rm \\Delta V_{FWHM}\\approx 5.1~km~s^{-1}$) and a plateau component ($\\rm \\Delta V_{FWHM}\\approx 13.4~km~s^{-1}$). The identities of numerous  U lines  may be either from unknown  transitions of the known molecules with complex rotational structure or from any transitions of unknown molecules. We leave the identifications of these lines to be explored in a future study.  \\subsection{Regional Origin of the Molecular Species} The spectral lines of each molecular species can often trace physically different regions of molecular clouds,  and so are useful for establishing the origin of the species and understanding the astro-chemical processes inside molecular clouds. Recently, Sutton et al (1995) have surveyed 5 distinct regions of the Orion molecular clouds with high spatial resolution (about $14''$ beam ) and characterized the kinematic properties and main tracers of the four distinguished spatial and kinematical components such as extended ridge, hot core, compact ridge, and plateau. Their data suggest that the extended ridge shows $\\rm V_{LSR} \\approx 9.8~km~s^{-1}$ and $\\rm \\Delta V_{FWHM} \\approx 2.5~km~s^{-1}$, and the hot core $\\rm V_{LSR} \\approx 7.0~km~s^{-1}$ and $\\rm \\Delta V_{FWHM} \\approx 8.0~km~s^{-1}$, the compact ridge $\\rm V_{LSR} \\approx 6.6~km~s^{-1}$ and $\\rm \\Delta V_{FWHM} \\approx 8.0~km~s^{-1}$, and the plateau $\\rm V_{LSR} \\approx 5.5\\sim 9.8~km~s^{-1}$ and $\\rm \\Delta V_{FWHM} \\approx 10 \\sim 20~km~s^{-1}$. The beam of the TRAO telescope at 140 GHz is more than 3 times bigger than the JCMT beam at 330 GHz, and so  encompasses nearly all components by Sutton et al. (1995). By comparing their kinematical values of these distinct components in the Orion clouds with our data, we  are able to determine which component a given molecular species traces. For this purpose, we obtained the $\\rm V_{LSR}$ and $\\rm \\Delta V_{FWHM}$ of each detected transition from all species from Gaussian fits to the spectral lines (Table 11), which are also plotted in Fig. 4. The values for the molecules with multiple transitions are given in average. According to the kinematic values of the species, $\\rm CH_2CHCN$ is thought to be mainly from the extended ridge and $\\rm C_2H_5CN$, $\\rm CH_3CN$, $\\rm H_2CO$, $\\rm SO_2$, $\\rm NO$, $\\rm CS$, $\\rm HC_3N$, and $\\rm OCS$ from the plateau. The kinematic values of the remaining species are similar to the values of the hot core or the compact ridge, and so it is difficult to discriminate where they are from, between the two regions."
    },
    "0011/astro-ph0011012_arXiv.txt": {
        "abstract": "Centrifugal acceleration of charged test particles at the base of a rotating jet magnetosphere is considered. Based on an analysis of forces we derive the equation for the radial accelerated motion and present an analytical solution. It is shown that for particles moving outwards along rotating magnetic field lines, the energy gain is in particular limited by the breakdown of the bead-on-the-wire approximation which occurs in the vicinity of the light cylinder $r_{\\rm L}$. The corresponding upper limit for the maximum Lorentz factor $\\gamma_{\\rm max}$ for electrons scales $\\propto B^{2/3} \\,r_{\\rm L}^{2/3}$, with $B$ the magnetic field strength at $r_{\\rm L}$, and is at most of the order of a $10^2-10^3$ for the conditions regarded to be typical for BL~Lac objects. Such values suggest that this mechanism may provide pre-accelerated seed particles which are required for efficient Fermi-type particle acceleration at larger scales in radio jets. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011538_arXiv.txt": {
        "abstract": "We have developed a novel computer code designed to follow the evolution of cosmic-ray modified shocks, including the full momentum dependence of the particles for a realistic diffusion coefficient model. In this form the problem is technically very difficult, because one needs to cover a wide range of diffusive scales, beginning with those slightly larger than the physical shock thickness. With most finite difference schemes for Euler's equations the numerical shock thickness is at least one zone across, so this provides a lower bound on the physical scale for diffusive transport computation. Our code uses sub-zone shock tracking \\citep{levshy95} and multi-level adaptive mesh refinement \\citep{berglev98} to provide enhanced spatial resolution around shocks at modest cost compared to the coarse grid and vastly improved cost effectiveness compared to a uniform, highly refined grid. We present and discuss the implications from our initial results. ",
        "introduction": "Diffusive shock acceleration (DSA) is now widely accepted as the model to explain the production of cosmic rays (CR) in a wide range of astrophysical environments \\citep{dru83,blaeic87,berzkry88}. The concept behind DSA, first-order Fermi acceleration of charged particles trapped between convergent flows across a shock, is quite simple. However, the full DSA problem is actually extremely complex, because the nonlinear interactions between energetic particles, resonantly scattering waves and the underlying plasma can become dominant effects. Important consequences of nonlinear interactions include such things as generation and damping of the scattering wave field, injection of suprathermal particles into the CR population, as well as heating and compression of the plasma flow due to the CR pressure. Owing to these complex nonlinear physics involved in the model, numerical simulations have been quite useful and successful in understanding the details of the acceleration process and dynamical feedback of the CRs to the underlying plasma \\citep{falgid87,EMP90,dorfi90,kanjon91,berz94,berzvolk00}.  In continuum approaches to numerical simulations of DSA theory, the CR diffusion-convection equation is solved at each of a large number of suprathermal momentum values simultaneously with a set of fluid equations describing the flow associated with the bulk, thermal plasma, including the nonlinear interactions between the plasma, CRs and scattering waves. Particle acceleration is effected by diffusion across velocity gradients in the motion of the scattering centers, which are usually assumed to be tied to the bulk flow. Pressure by the diffusing CRs, in turn, decelerates and compresses flow into the shock, forming a shock ``precursor''. Since that development eliminates the original, simple velocity jump seen by the CRs, the DSA is then modified according to details of the flow within the precursor, whose scales are characterized by the so-called diffusion length, $D_{\\rm diff}(p) = \\kappa(p)/u$, where $\\kappa$ is the spatial diffusion coefficient for CRs of momentum $p$, and $u$ is the characteristic flow velocity against which the CRs must swim, \\eg \\citep{kanjon91}. Accurate solutions to the CR diffusion-convection equation require a computational grid spacing significantly smaller than $D_{\\rm diff}$, typically, $\\Delta x \\sim 0.05 D_{\\rm diff}(p)$. In a realistic diffusion transport model, it is thought that the diffusion coefficient should have a steep momentum dependence, $\\kappa (p) \\propto p^s$, with $s \\sim 1-2$. For the lowest energy CR particles the diffusion lengths ($D_{\\rm diff}(p)$) are only slightly greater than the shock thickness, while they can be many orders of magnitude greater than that for the highest energy particles. Thus, a wide range of length scales is required to be resolved in order to solve the diffusion convection equation correctly for the model with a realistic diffusion coefficient. Previous numerical simulations which adopted the traditional flux-differencing method on a uniform grid were often forced to assume a weak momentum dependence, for example,  $s=0.25$ in \\citet{kanjon91}.  To overcome this numerical problem, \\citet{berz94} introduced a ``change of variables technique'' in which the radial coordinate is transformed into a new variable, $x(p)= \\exp[-(r-R_s)/D_{\\rm diff}(p)]$ where $R_s$ is the shock radius, defined for each particle momentum for the upstream region. A uniform grid is used for the downstream region. This allowed them to solve the coupled system of gasdynamic equations and the CR transport equation even when the diffusion coefficient has a strong momentum dependence (\\eg $\\kappa(p)\\propto p$). Their code is designed for simulations of supernova remnants, which are represented by piston-driven spherical shocks in one-dimensional geometry. It is different from conventional Eulerian codes in several ways. Both gasdynamic equations and the CR transport equation are solved separately either side of the gas subshock. Then the gasdynamic solutions at both sides of the subshock are used to solve the Riemann problem, which determines how the subshock evolves. Also an iteration scheme is applied to match the downstream and upstream solutions for the CR diffusion-convection equation at the subshock. In any case this has enabled them to explore several important issues regarding the particle acceleration at supernova remnants more fully than was possible before, \\eg \\citet{berz95,berz96,berzvolk00}. However, no consistency-check for this method has been attempted so far, since no existing conventional codes can handle such a strongly momentum dependent diffusion coefficient.  Fermi shock acceleration affects those particles with a mean free path greater than the shock thickness that can resonantly scatter with self-generated Alfv\\'en waves. In the so-called ``thermal leakage'' type injection model, the diffusion and acceleration of these particles out of the suprathermal tail of the Maxwellian distribution determines the CR injection rate \\citep{elleich84,kanjon95}. A self--consistent, analytic and nonlinear model for ion injection based on the interactions of the suprathermal particles with self--generated magneto--hydrodynamic waves in strong shocks has been presented by \\citet{mal98}. By adopting this analytic solution, \\citet{gies00} have developed a numerical treatment of the injection model at a strong quasi--parallel shock, which is then incorporated into the combined gas dynamics and the CR diffusion--convection code. Since the suprathermal particles have mean free paths a few times that of thermal particles, resolving these smallest scales is of critical importance in estimating the injection and acceleration efficiency in such numerical simulations of the CR modified shocks. In fact, \\citet{gies00} were able to run their simulations, with a conventional Eulerian scheme on a uniform grid, only up to the time when the maximum accelerated momentum was of order of $p_{\\rm max}/m_{\\rm p}c \\sim 1$ for a Bohm type diffusion model because of severe requirements for computational resources needed to evolve the CR distribution to highly relativistic momenta. This calls for an alternative method comparable to Berezhko's code, which solves the CR diffusion-convection equation on a grid whose spacing scales with the diffusion length at each momentum value sampled.  In this contribution, we present a new numerical scheme that follows CR modified shocks in one dimensional, plane-parallel geometry. We take advantage of the fact that the diffusion and acceleration of the low energy particles are important only close to the shock owing to their small diffusion lengths. They are simply advected along with the underlying gas flow far upstream and downstream of the shock. Thus it is necessary to resolve numerically the diffusion length of the particles only around the shock. So we first implement a shock tracking scheme to locate the shock position exactly and then increase the grid resolution only around the shock by applying multi-levels of refined grids. Toward this end, we have adopted the shock tracking method of \\citet{levshy95} and the Adaptive Mesh Refinement (AMR) technique of \\citet{berglev98}, and modified the code to use multiple levels of grid refinement only around the shock.  In the following section we outline our numerical methods, while in \\S 3 we present and discuss our test results. Section 4 provides a summary. ",
        "conclusions": "We have developed a new hydro/CR dynamics code which is specifically designed to solve the time dependent evolution of CR shocks. Diffusive shock acceleration of the CR particles depends on the diffusion of particles whose momenta span many orders of magnitude. Since the length and time scales for evolution of the CR kinetic equation scale directly with the diffusion coefficient, an accurate solution to the problem requires that one include all of those scales in the simulation, beginning just outside the gas subshock thickness. Thus, in order to follow accurately the evolution of a CR modified shock, it is necessary to resolve the precursor structure upstream of the subshock and, at the same time, to solve correctly the diffusion of the low energy particles near the injection pool. These low energy particles have diffusion lengths that are much smaller than the scale length of the precursor, so a large dynamic range of resolved scales is required for CR shock simulations. To solve this problem generally we have successfully combined a powerful ``Adaptive Mesh Refinement'' (AMR) technique \\citep{bergcol89,berglev98} and a ``shock tracking'' technique \\citep{levshy95}, and implemented them into a hydro/CR code based on the wave-propagation method \\citep{lev97}. The AMR technique allows us to ``zoom in'' inside the precursor structure with a hierarchy of small, refined grid levels applied around the shock. The shock tracking technique tracks hydrodynamical shocks and maintains them as true discontinuities, thus allowing us to refine the region around the shock at an arbitrary level. The result is an enormous savings in both computational time and data storage over what would be required to solve the problem using more traditional methods on a single fine grid.  The code has been applied to simulations of CR modified shocks with several diffusion coefficient models with strong momentum dependence, which were not possible previously due to severe computational requirements. The main conclusions from the simulations can be summarized as follows:  \\begin{enumerate}  \\item{} Our CR/AMR technique code proves to be very cost effective. In typical simulations where 10\\% of the base grid is refined with $l_{\\rm max}$ levels, for example, the computing time increases by factors of $(2^{l_{\\rm max}})^{0.7}$ compared with the case of no refinement ($l_{\\rm max}=0$). It should be compared with the time increases by factors of $ (2^{l_{\\rm max}})^2$ for the simulations of an uniform grid spacing that matches the cell size at the $l_{\\rm max}-th$ refined level grid. In a simulation where the precursor scale is only a small fraction of the computational domain, the advantage in computing time of the refined multi-level grid over a single fine grid becomes even greater.  \\item{} A convergence test is performed for a Mach 20 gas shock with $V_s/c=0.01$, which evolves into a CR pressure dominated shock. Comparison between a simulation on a coarse grid with multi-level refinement and another simulation on a single fine grid without refinement has demonstrated that our CR/AMR code can perform the intended simulation with reasonable accuracy at a much shorter computing time. The difference in the CR pressure in two test simulations is around 10\\% in early evolution, but deceases to a few \\% after the shock has reached a quasi-steady state in later evolution. The required computing time is reduced by a factor of 150 in the AMR simulation.  \\item{} We also carried out a set of simulations when three different diffusion models, $\\kappa_{\\rm l}= p/\\sqrt{2}$, $p^{1.5}/\\sqrt{2}$, and $p^2/(p^2+1)^{1/2}$ for $p<1$, are included, while a Bohm type, $\\kappa_{\\rm h}= p^2/(p^2+1)^{1/2}$ is assumed for $p\\ge1$. Three simulations generate similar results once the CR pressure is dominated by the relativistic particles ($p>1$), when the maximum acceleration momentum becomes $p_{\\rm max} >>1$.  Thus a diffusion model of $\\kappa \\propto p$ can be used instead of a Bohm model as long as one does not focus on the early evolution when $p_{\\rm max}<1$. Since we can use much larger grid spacings in simulations with $\\kappa \\propto p$ model than those with a Bohm type diffusion at low momenta $p<<1$, the required level of refinements and so the computing resources can be reduced greatly.  \\item{} For a Mach 20 shock, with an injection rate of $\\eta \\sim 6\\times 10^{-4}$, the shock becomes CR dominated and develops a significant precursor. Since the flow is decelerated gradually through the precursor, the velocity jump that the CR particles sample across the shock depends on the diffusion length of the particle, that is, $\\Delta u(p) = fcn (\\kappa(p))$. So the slope of the particle distribution function, defined as $q(p)=-(\\partial \\ln{f}/\\partial\\ln{p})$, increases with $p$. In the simulated shock, the compression ratios across the subshock and across the total transition are 3.1 and 11, respectively, so $f(p)$ is $\\propto p^{-4.5}$ at low energy momenta but flattens to $ f(p) \\propto p^{-3.3}$ at high energy momenta just below $p_{\\rm max}$. This demonstrates that nonlinear feedbacks between the precursor dynamics and the CR injection and acceleration should be treated accurately in numerical simulations of CR shocks.  \\end{enumerate}  We are currently implementing a numerical scheme for the self-consistent injection model by \\citet{gies00} which is based on the plasma-physics study of the nonlinear wave-particle interactions in shocks presented by \\citet{mal98}. This will allow us to eliminate any free parameters for the injection process from the CR shock simulations. We intend also to extend the code to treat spherical shocks in order to study CR acceleration in supernova remnants."
    },
    "0011/astro-ph0011224_arXiv.txt": {
        "abstract": "Star formation is intimately linked to the dynamical evolution of molecular clouds. Turbulent fragmentation determines where and when protostellar cores form, and how they contract and grow in mass via accretion from the surrounding cloud material. Efficiency, spatial distribution and timescale of star formation in turbulent clouds are estimated by comparing numerical models of self-gravitating isothermal gas where turbulence is assumed to have decayed or is driven at supersonic speed on various wavelengths.  Turbulence that is not continuously replenished or that is driven on large scales leads to a rapid formation of stars in a clustered mode, whereas interstellar turbulence that carries most energy on small scales results in isolated star formation with low efficiency. Protostellar accretion rates strongly vary in time and peak at a few $10^{-5}\\,$M$_{\\odot}\\,$yr$^{-1}$. The clump mass spectrum for models of pure hydrodynamic turbulence is steeper than the observed one, but becomes comparable when including gravity. The mass spectrum of dense cores, on the other hand, is log-normal for decaying and large-wavelength turbulence, similar to the IMF, but is to flat in the case of small-scale turbulence. ",
        "introduction": "Understanding the processes that lead to the formation of stars is one of the fundamental challenges in astronomy.  Stars are born in turbulent interstellar clouds of molecular hydrogen. The location and the mass growth of young stars are hereby intimately coupled to the dynamical cloud environment. Stars form by gravitational collapse of shock compressed density fluctuation generated from the supersonic turbulence ubiquitously observed in molecular cloud (e.g.\\ Padoan 1995, Klessen, Heitsch, \\& Mac~Low 2000). Once a gas clump becomes gravitationally unstable, collapse progresses in the center until a protostellar object forms which continues to grow in mass via accretion from the infalling envelope (see Wuchterl, this volume). Also, stars hardly ever form in isolation. Instead they come in small aggregates or larger clusters (Lada 1992, Mizuno et al.\\ 1995, Testi, Palla, \\& Natta 1998), where the interaction of protostellar cores and their competition for mass from their surrounding are important processes shaping the distribution of the final star properties.   Star formation can thus be seen as a two-phase process: First, {\\em turbulent fragmentation} leads to transient clumpy molecular cloud structure, with some of the density fluctuation exceeding the critical mass and density for gravitational contraction. This sets the stage for the second phase of star formation, the {\\em collapse of individual protostellar clumps} building up central protostars. In the following, the words gas clump and shock-generated density fluctuation are used synonymously. Cores are defined as the high-density central regions of collapsing clumps where protostars build up. I focus  on the first phase of star formation and demonstrate how the turbulent interstellar velocity field influences the properties of forming stellar aggregates and clusters. Using numerical simulations (\\S{2}) of decayed and driven supersonic turbulence in self-gravitating isothermal gas as templates for molecular cloud dynamics, I discuss the spatial distribution and timescale (\\S{3}) of star formation, the resulting protostellar accretion rates (\\S{4}), and the expected mass spectra for gas clumps and protostellar cores (\\S{5}). ",
        "conclusions": ""
    },
    "0011/astro-ph0011068_arXiv.txt": {
        "abstract": "This work provides an in-depth mathematical description of the response functions that are used for spatial and spectral analysis of X-ray data.  The use of such functions is well-known to anyone familiar with the analysis of X-ray data where they may be identified with the quantities contained in the Ancillary Response File (ARF), the Redistribution Matrix File (RMF), and the Exposure Map.  Starting from first-principles, explicit mathematical expressions for these functions, for both imaging and dispersive modes, are arrived at in terms of the underlying instrumental characteristics of the telescope including the effects of pointing motion.  The response functions are presented in the context of integral equations relating the expected detector count rate to the source spectrum incident upon the telescope.  Their application to the analysis of several source distributions is considered.  These include multiple, possibly overlapping, and spectrally distinct point sources, as well as extended sources.  Assumptions and limitations behind the usage of these functions, as well as their practical computation are addressed. ",
        "introduction": "%  It is a basic tenet of X-ray spectral analysis that the source flux incident at the telescope is related to the observed count rate through an integral equation involving the effective area of the telescope.  The most commonly accepted technique for dealing with this equation involves the use of spectral analysis programs such as \\verb|xspec| \\citep{xspecmanual}.  The effective area is input into these programs via a file called the Ancillary Response File, or ARF. In addition, the energy resolution of the detector is specified by the Redistribution Matrix File, or RMF. This work presents formal descriptions of the quantities embodied by the ARF and RMF in terms of the underlying instrumental responses, making a clear connection between the incident source flux and the observed count rate.  Roughly speaking, the effective area of an X-ray telescope composed of a mirror and a detector is more or less the product of the effective area of the mirror with the quantum efficiency (QE) of the detector. However, the ARF, which relates observed counts to a source flux, also depends upon observation-dependent quantities such as the detailed aspect history of the telescope, its point spread function (PSF), and upon details of the analysis itself, e.g., the filtering and binning of the observed data.  With the advent of the Chandra X-ray Observatory \\citep{chandra}, all of the subtleties introduced by the shape of the PSF and the telescope aspect motion were deemed important for the computation of an ARF. Chandra is well calibrated \\citep{chandracal} and, with its unprecedented combined spectral and spatial response, a precise definition of the ARF that incorporates the effects of spacecraft motion and the PSF is necessary in order to perform spectral analysis at the highest resolution of the instrument.  The main goal of this work is to present an explicit first-principles derivation of the ARF by including the proper treatment of telescope motion (e.g., dither) and the PSF, as well as any filtering of the data.  The bottom-up approach taken here necessarily implies that a meaningful and consistent derivation can be achieved by considering the role of the ARF in spectral analysis.  As a result, the ARF is presented in the context of integral equations that connect the incident X-ray flux to an expected count rate.  Traditionally, the ARF and RMF have been used primarily for the analysis of spectral image data.  An important aspect of this work is to extend this approach to the analysis of dispersive spectral data, such as data obtained by Chandra or Newton. To this end, definitions of an ARF and an RMF are presented that are suitable for dispersive data analysis, and which may be utilized by existing spectral analysis software.  One of the original motivations for this work was the need to create a related object, an exposure map, for use in the analysis of data obtained by the Chandra X-ray Observatory.  This paper also gives a rigorous definition of the exposure map and discusses some of its uses and its limitations in spectral image analysis.  The resulting definition is consistent with current use, intuition, and physics.  The next section contains a discussion of the general response of an X-ray telescope and also serves to introduce the notation and conventions used throughout this paper.  Although originally inspired by the need to create ARFs and exposure maps for Chandra, the presentation has been kept as general as possible without focusing on any particular telescope or instrument. A derivation of the imaging ARF follows in section \\ref{sec:arfdef} where its application to several problems is considered.  These include the problem of multiple overlapping point sources.  Section \\ref{sec:exposuremap} contains a definition of the exposure map and explores its use as well as its limitations in dealing with extended sources.  The definition of a dispersive ARF and RMF that are suitable for use in the analysis of dispersive spectral data are given in section \\ref{sec:garf}. Following the summary of the paper is an appendix that considers the practical computation of these objects. ",
        "conclusions": "In this paper, explicit expressions for the imaging ARF, grating ARF, and the exposure map were given in terms of the underlying instrumental responses that are consistent with the current use of the objects.  These quantities were obtained from first principles by relating the expected detector count-rate to an incident photon source flux via the overall telescope response function suitably factored into individual instrumental responses.  One of the complications in the derivation of these quantities concerned the proper treatment of time-varying effects due to telescope pointing motion, e.g., dither.  At the same time, the assumed presence of motion about some nominal pointing allowed some important factorizations to take place that otherwise would have been suspect in regions containing detector boundaries or bad pixels.  For this reason, purposely dithering an observation is recommended, provided, of course, that one can reconstruct the aspect history with sufficient accuracy.  An added benefit of the first principles approach taken here is that it allows one to consider problems that cannot readily be handled by conventional means through the use of an ARF and an RMF. For example, as shown in section \\ref{sec:extended}, it ARF may be applied to the analysis of an extended source provided one knows {\\em a priori} that the source flux distribution factors into a known spatial component and an unknown spectral component.  It is easy to find sources where such a factorization is not permissible; the supernova remanent, Cassiopeia A, is one.  Another problem that does not appear to be treatable through standard the techniques is the analysis of an extended source in the presence of a diffraction grating. The grating ARF defined in section \\ref{sec:garf} was derived assuming a point source distribution.  The basic problem with the analysis of an extended source is that, unlike a point source, there is no unique zeroth order position that one could use in the grating equation.  By judicious filtering in pulse-height space, one may find regions where there is enough of a point-like behavior to permit the grating ARF to be used. However, how to handle a generic extended source in the presence of a diffraction grating is still an open question.  It is hoped that the mathematical formulation of the extended source problem as given in sections \\ref{sec:arfdef} and \\ref{sec:garf} will lead to better insights into these problems and ultimately to their solution.  This work also highlights some important practical considerations that should be taken into account in the design of astronomical data analysis software systems.  For instance, to allow for the possibility of spatial variation in the underlying detector redistribution function, the software component responsible for the filtering of events should allow the user to easily filter {\\em simultaneously} on both sky coordinates and detector coordinates.  In addition, both filters would need to be passed to the program that generates the ARF. Finally, spectral fitting programs should be enhanced to facilitate the analysis of blended sources by handling the coupled integral equations in \\eq{intsys}."
    },
    "0011/astro-ph0011542_arXiv.txt": {
        "abstract": "We present RXTE monitoring observations of the Seyfert 2 galaxy Mrk 348 spanning a 6 month period. The time-averaged spectrum in the 3--20 keV band shows many features characteristic of a Compton-thin Seyfert 2 galaxy, namely a hard underlying power-law continuum ($\\Gamma \\approx 1.8$) with heavy soft X-ray absorption ($N_{H} \\sim 10^{23}$~cm$^{-2}$) plus measureable iron K$\\alpha$ emission (equivalent width $\\sim 100$ eV) and, at high energy, evidence for a reflection component ($R \\simless 1$). During the first half of the monitoring period the X-ray continuum flux from Mrk 348 remained relatively steady. However this was followed by a significant brightening of the source (by roughly a factor of 4) with the fastest change corresponding to a doubling of its X-ray flux on a timescale of about 20 days. The flux increase was accompanied by a marked softening of X-ray spectrum most likely attributable to a factor $\\sim 3$ decline in the intrinsic line-of-sight column density.  In contrast the iron K$\\alpha$ line and the reflection components showed no evidence of variability.  These observations suggest a scenario in which the central X-ray source is surrounded by a patchy distribution of absorbing material located within about a light-week of the nucleus of Mrk 348.  The random movement of individual clouds within the absorbing screen, across our line of sight, produces substantial temporal variations in the measured column density on timescales of weeks to months and gives rise to the observed X-ray spectral variability. However, as viewed from the nucleus the global coverage and typical thickness of the cloud layer remains relatively constant. ",
        "introduction": "\\label{sec-intro}  The discovery of ``hidden'' broad-line regions (BLRs) in several Seyfert 2 galaxies \\citep{am85,mg90,tra92} has generated much interest in unified theories, in which the viewing orientation explains many of the observed differences between type 1 and 2 Seyfert galaxies (Antonucci 1993 and references therein). Such theories posit that, in both types of Seyfert galaxy, a luminous central source (presumably an accreting supermassive blackhole) is surrounded by a equatorial distribution of matter in the form of either a swollen accretion disk or a molecular torus \\citep{kb86}. A type 1 nucleus is then observed only if the viewing direction is sufficiently close to the axis of the system so as to fall within the opening angle of the putative torus. More inclined lines of sight necessarily intercept the disk/torus material in which case we see an ``obscured'' type 2 system.  However, a caveat is that a small fraction of the nuclear flux can be scattered into our line of sight by highly ionized material filling the ``hole'' of the torus, thus explaining the detection of broad optical lines in polarized light and also weak UV and soft X-ray emission even in Seyfert 2s where our direct view of the nucleus is completely blocked.  X-ray measurements of intrinsic column densities of $\\simgreat 10^{22}$ cm$^{-2}$ in many Seyfert 2 galaxies \\citep{awa91,sd96,tur97a} lend support to the above unification picture. For column densities $\\simless 10^{24}$ cm$^{-2}$, X-rays can penetrate the torus leading to the detection of an absorbed power-law continuum together with a iron K$\\alpha$ fluorescence line at 6.4~keV. However for column densities $\\simgreat 10^{24}$ cm$^{-2}$, even hard X-radiation is seen only indirectly via scattering (as in Compton-thick Seyfert 2s such as NGC 1068).  X-rays may be electron scattered by highly ionized matter extending along the axis of the system or Compton reflected by almost neutral material located at the inner edge of the torus.  In either case, the scattering material is exposed to the intense nuclear flux and may produce significant line emission and, in circumstances where the direct continuum is completely blocked leaving only a baseline of the much weaker scattered continuum, such lines can have very high equivalent widths \\citep{ghi94,kro94}.  The spectral indices found in Seyfert galaxies are distributed around a single value of $\\Gamma \\approx 1.9$ \\citep{np94,sd96}.  The best explanation for this is that mildly relativistic thermal electrons or highly relativistic non-thermal electrons Compton scatter UV photons into the X-ray energy range (e.g. Svensson 1996 and references therein).  The rollover seen in X-ray data above 100 keV has focussed most attention on thermal models \\citep{zdz95,gon96}.  A popular scenario is the two-phase disk-corona model, in which a hot X-ray corona is located above a cold UV emitting accretion disk \\citep{hm91,hm93}.  The UV seed photons for Compton cooling of the energetic electrons are produced from reprocessing of the hard X-ray spectrum in the accretion disk.  With complete feedback, the approximate equipartition between the soft disk and hard X-ray luminosities leads naturally to $\\Gamma \\simgreat 1.9$ for a wide range of optical depths (see also Stern et al. 1995).  However, for sources with spectra flatter than $\\Gamma \\approx 1.9$, the disk-corona model of \\citet{hm91,hm93} must be modified so that some of the soft seed photons escape without being intercepted by the corona, thus making for harder spectra as fewer photons are available for Compton scattering \\citep{haa94}.  X-ray variability studies of Seyfert 2 nuclei can provide important clues to the processes by which X-rays are produced as well as to the geometry of the circumnuclear matter.  However, variability has been reported in only a few objects.  One example is Mrk 3, where the hard X-ray emission decreased by a factor of two during a period of 3.6 years \\citep{awa90,mar92,iwa94}.  Recent observations with \\emph{RXTE} have shown variability to exist on time-scales of weeks \\citep{geo99}. The shape of the spectrum suggests that \\emph{all} of the variability can be attributed to changes in the direct nuclear flux \\citep{cap99}. Short term variability (on time-scales of hours) has been detected in only one object, NGC 4945 \\citep{iwa93,gua00}.  These observations suggest that a parsec scale, geometrically thick molecular torus cannot be responsible for the bulk of the X-ray absorption \\citep{mad00}.  Spectral variability has also been claimed in a couple of Seyfert 2 galaxies.  For example, the absorption column in NGC 7582 increased by $N_{\\rm H} \\sim 4 \\times 10^{22}$ cm$^{-2}$ between \\emph{ASCA} observations taken two years apart \\citep{xue98}.  A similar variation in the absorber was suggested by \\citet{war93}, based on a comparison between \\emph{Einstein}, \\emph{EXOSAT}, and \\emph{Ginga} observations. More recent observations with \\emph{BeppoSAX} of this particular galaxy have shown that the absorber may have a complex spatial structure with a large column density ($N_{\\rm H} \\sim 10^{24}$ cm$^{-2}$) covering 60\\% of hard X-ray continuum, and a smaller, but variable column completely covering the source \\citep{tur00}.  As a further example, in NGC 7172, the spectral index of the power law continuum decreased by $\\Delta \\Gamma \\approx 0.3$ between \\emph{Ginga} and \\emph{ASCA} observations taken almost 4 years apart \\citep{ryd97}.  These authors attribute the decrease to an intrinsic change in the spectral index although, given the narrow bandpass of \\emph{ASCA}, it is possible that the spectrum is affected by a complex absorber and/or Compton reflection.  The launch of the \\emph{RXTE} satellite is conducive to the systematic study of variability in Seyfert 2 galaxies, because of the large effective area and broad band coverage afforded by the detectors, and the flexible scheduling allowed by the mission.  We have conducted a six month monitoring campaign with \\emph{RXTE} of a few Seyfert 2 galaxies previously observed by the \\emph{Ginga} and \\emph{ASCA} satellites.  The results of our observations of Mrk 3 have been reported elsewhere \\citep{geo99}.  Here we present our results on Mrk 348.  Mrk~348 is a nearby ($z=0.015$; de Vaucouleurs et al. 1991) Seyfert 2 galaxy for which there is evidence of a broad (FWHM $\\sim 7400$ km~s$^{-1}$) H$\\alpha$ line component in polarized light \\citep{mg90}. In terms of current unification schemes for Seyfert galaxies (see above) this can be taken as clear evidence for the presence of an obscured Seyfert 1 nucleus, confirmation of which was provided by the \\emph{Ginga} detection of an absorbed ($N_{\\rm H} = 10^{23.10\\pm0.08}$ cm$^{-2}$) hard X-ray source (with $\\Gamma = 1.68 \\pm 0.17$) coincident with this galaxy \\citep{war89}.  This heavy obscuration in the nucleus of Mrk 348 may also be responsible for the bi-canonical nature of its ionizing continuum \\citep{mul96,sim96}.  There was marginal evidence in the \\emph{Ginga} spectrum for an iron K$\\alpha$ emission line at 6.4 keV of equivalent width $130\\pm130$ eV. An iron K$\\alpha$ emission line of equivalent width $113\\pm56$ eV was subsequently detected with \\emph{ASCA} \\citep{net98}.  In addition to iron K$\\alpha$, \\emph{ASCA} detected a prominent S {\\sc xiv}--{\\sc xv\\/} line of large equivalent width ($600\\pm300$ eV).  A weak power law continuum was detected at soft X-ray energies in observations made with the \\emph{ROSAT} Position Sensitivity Proportional Counter (PSPC) ($\\Gamma = 2.40^{+0.58}_{-0.42}$; Mulchaey et al. 1993) and \\emph{ASCA} ($\\Gamma = 2.77\\pm0.87$; Netzer et al. 1998).  It is most likely that this weak soft X-ray emission is dominated by electron-scattered nuclear flux, with perhaps some contribution from spatially extended thermal emission from a hot wind (e.g. Krolik \\& Vrtilek 1984).  The 2--10 keV luminosity measured in the \\emph{Ginga} observation was $1.2 \\times 10^{43}$ erg s$^{-1}$ ($H_{\\circ} = 50; q_{\\circ} = 0.5$), which is roughly a factor 3 higher than that subsequently recorded by \\emph{ASCA}. ",
        "conclusions": "\\label{sec-discussion}  The shortest timescale on which X-ray continuum variations are observed in Mrk 348 is $\\sim 1$~day, the minimum spacing between observations.  In terms of the light-travel time this implies a maximum source size of $r_{\\rm s} \\simeq 2.5 \\times 10^{15}$~cm, although in reality the X-ray emitting region may be much smaller than this, (i.e. $\\sim 3 \\times 10^{14}$~cm if the X-rays are produced within 10~Schwarzchild radii of a $10^{8}~M_{\\odot}$ black hole). In contrast variations in the intrinsic column density (along the line of sight to the nucleus of Mrk 348) take place over periods of typically weeks to months with the largest change, $\\Delta N_{\\rm H} \\sim 10^{23}$~cm$^{-2}$, occurring on a timescale of $\\sim 70$ days.  If the observed column changes are interpreted as due to the motion, across the line of sight, of clouds of gaseous material in orbit around the central black-hole, then for a transverse cloud velocity of $v_{\\rm cloud} \\approx 10,000$~km~s$^{-1}$ (characteristic of extreme BLR cloud velocities), the predicted timescale for the column change is $r_{\\rm s}/v_{\\rm cloud} \\simeq 30$~days, which is roughly consistent with the observations.  Such clouds are located at a distance of $R = GM/v_{\\rm cloud}^{2} \\simeq 1.3 \\times 10^{16} M_{8}$~cm, where $M_{8}$ is the black-hole mass in units of $10^{8}~M_{\\odot}$.  We can infer the ionization of the cloud material via the ionization parameter $\\xi$, where $\\xi = L_{\\rm ion}/n_{\\rm e} R^{2}$.  Here $L_{\\rm ion}$ is the ionizing luminosity (measured between 13.6~eV to 13.6~keV; Kallman \\& McCray 1982), $n_{\\rm e}$ is the gas density and $R$ is the distance to the clouds from the continuum source.  For an ionizing luminosity of $\\sim 5 \\times 10^{43}$~erg~s$^{-1}$ and $n_{\\rm e} \\approx 10^{10}$~cm$^{-3}$ (a typical value for BLR clouds), the ionization parameter is $\\xi \\approx 30$. This is consistent with a lightly ionized state for the iron in the clouds (i.e. below Fe{\\sc xv}), but in agreement with the iron K$\\alpha$ emission-line (and K-edge) requirements.  However for a typical cloud column density of $\\Delta N_{\\rm H} \\sim 10^{23}$~cm$^{-2}$, the implied physical thickness, is only $\\Delta N_{\\rm H}/n_{\\rm e} \\sim 10^{13}$ cm. Since this dimension is much smaller than the inferred X-ray source size $r_{\\rm s}$, the organisation of clouds into a coherent sheet-like distribution is required in order to explain the large amplitude $N_{\\rm H}$ variations which are observed in Mrk 348 and other sources (e.g. Yaqoob et al. 1993).  Alternatively, the increase in the source luminosity may be sufficient to induce the observed decline in the column density via photo-ionization.  However, the lack of evidence for highly ionized iron in Mrk 348 would require a fairly narrow regime for the ionization parameter.  Such a scenario could be investigated using photo-ionization codes such as {\\sc xstar\\/}, but this is beyond the scope of the present paper.  Within the measurement constraints, the iron K$\\alpha$ line flux remained at a roughly constant level throughout the monitoring campaign.  This implies that much of the line emission is produced in material reasonably distant from the nucleus.  A possible location for this material is in the patchy cloud distribution responsible for the highly variable intrinsic absorption in this source.  The lack of significant changes in the line flux then requires both the sky coverage and typical thickness of the cloud layer, as viewed from the nucleus, to remain fairly constant (a requirement which need not, of course, conflict with the fact that temporal variations are observed along our specific line of sight).  It is also possible that a contribution to the line emission comes from the illumination of the (far-side) inner walls of the putative molecular torus (assuming at least some element of the gaseous medium surrounding the X-ray source is Thomson thick).  While Compton reflection is not specifically required by the spectral fitting, the data are consistent with the reflection expected from a torus subtending a solid angle at the source of $\\Omega/2\\pi \\approx 0.3$--$0.8$.  If the ionization cones are collimated by the torus, then the observed half-opening angle of the cones, of $\\theta \\sim 45^{\\circ}$ \\citep{sim96}, corresponds to a torus solid angle of order $\\Omega/2\\pi = 2 \\cos \\theta = 1.4$.  Thus, in principle the bulk of the line emission could originate in this way. In this setting, the fact that Mrk 348 is not observed as a Compton-thick Seyfert 2 is presumably the result of our line of sight to the nucleus passing across the top of the torus thus avoiding its most optically thick regions.  Finally, we note that the iron K$\\alpha$ emission observed in the X-ray spectra of many Seyfert~1 galaxies is most naturally associated with reflection from the inner regions of an accretion disk.  One signature of such emission is the broadening of the line due to relativistic effects close to the central massive black-hole (e.g. Tanaka et al. 1995).  Also, depending on the geometry and optical depth of the X-ray emitting plasma, variations in the line flux on timescales of a week or less are predicted. The current observations provide no evidence for such an origin for the iron K$\\alpha$ emission observed in Mrk 348.  Also an earlier \\emph{ASCA} observation of Mrk 348 was unable to determine whether the iron K$\\alpha$ emission line was significantly broadened \\citep{net98}. Clearly further observations with much higher signal-to-noise and energy resolution will be needed in order to test for the presence of an accretion disk in Mrk 348."
    },
    "0011/astro-ph0011404_arXiv.txt": {
        "abstract": "Using X-ray data from the Rossi X-Ray Timing Explorer (RXTE), we report the pulse timing results of the accretion powered high mass X-ray binary (HMXRB) pulsar 4U 1907+09 covering a time span of almost two years. We measured three new pulse periods in addition to the previously measured four pulse periods. We are able to connect pulse arrival times in phase for more than a year. The source has been spinning down almost at a constant rate with a spin down rate of $\\dot \\nu$ = (-3.54 $\\pm  0.02) \\times 10^{-14}$ Hz s$^{-1}$ for more than 15 years. Residuals of pulse arrival times yield a very low level of random walk noise strengths $\\sim 2 \\times 10^{-20}$ rad$^{2}$ sec$^{-3}$ on a time scale of 383 days, which is four decades lower than that of the HMXRB pulsar Vela X-1. The noise strength is only a factor of 5 greater than that of the low mass X-ray binary pulsar (LMXRB) 4U 1626-67. The low level of the timing noise and the very stable spin down rate of 4U 1907+09 makes this source unique among the HMXRBs, providing another example, in addition to 4U 1626-67, of long term quiet spin down from an accreting source. These examples show that the extended quiet spin down episodes observed in the anomalous X-ray pulsars (AXPs) pulsars 1RXS J170849.0-400910 and 1E 2259+586 do not necessarly imply that these sources are not accreting pulsars. ",
        "introduction": "4U 1907+09 is an accretion powered X-ray binary pulsar which is accreting plasma from a blue supergiant companion star. It was discovered as an X-ray source by Giacconi et al. (1971) and has been studied using instruments on board  Ariel V (Marshall $\\&$ Ricketts 1980), Tenma (Makishima et al., 1984), EXOSAT (Cook $\\&$ Page 1987), Ginga (Makishima $\\&$ Mihara 1992, Mihara 1995), and RXTE (In 't Zand, Strohmayer $\\&$ Baykal 1997, In 't Zand, Baykal $\\&$ Strohmayer 1998, In 't Zand, Strohmayer $\\&$ Baykal 1998). Marshall $\\&$ Ricketts (1980) first determined the orbital period of the binary at 8.38 days by analysing the data taken between 1974 and 1980 from a survey instrument on board Ariel V. Folding the light curve of these data, they found two flares, a primary and a secondary, each occurring at the same orbital phase. Subsequent Tenma observations of this source have shown a pulse period at 437.5 sec (Maksihima et al., 1984). Later EXOSAT (Cook $\\&$ Page 1987) and recent RXTE obsevations (In 't Zand, Baykal $\\&$ Strohmayer 1998, In 't Zand, Strohmayer $\\&$ Baykal 1998) have shown that these flares are locked to orbital phases separated by half an orbital period. Makishima et al., (1984) and Cook $\\&$ Page (1987) suggested that the two flares are due to an equatorial disk-like envelope around a companion star which is inclined with respect to the orbital plane. When the neutron star crosses the disk, the mass accretion rate onto the neutron star, therefore the  X-ray flux, increases temporarily. Transient $\\sim 18$ sec oscillations have appeared during the secondary flare (In 't Zand, Baykal $\\&$ Strohmayer 1998). These oscillations may be interpreted as Keplerian motion of an accretion disk near the magnetospheric radius. Due to the long spin period the co-rotation radius is much larger than the magnetospheric radius corresponding to the magnetic field of 2.1 $\\times$ 10$^{12}$ Gauss implied by a cyclotron feature in the X-ray spectrum (Cusumano et al., 1998). Therefore 4U 1907+09 is not likely to be spinning near equilibrium, unlike other accretion powered X-ray pulsars. The 18 second quasi periodic oscillation at the flare suggests the formation of transient accretion disks from the wind accretion (In 't Zand, Baykal $\\&$ Strohmayer 1998). Another interesting feature of the source is the sudden decrease in X-ray intensity by a factor of $\\sim$ 10$^{2}$, during time intervals ranging from a few minutes to $\\sim$ 1.5 hour (In 't Zand, Strohmayer $\\&$ Baykal 1997). The spectra at the dipping activity and outside the dip periods were similar with no indication of large changes in the column density of cold circumstellar matter (i.e. N$_{H}$ remains below 10$^{23}$ cm$^{-2}$). It is suggested that the mass accretion rate ceased due to the inhomogeneous spherical wind from the companion.   In this work, we have investigated the stability of the spin down rate. This source has shown spin down rate changes less than $\\sim$ 8 $\\%$ within 12 years (In 't Zand, Strohmayer $\\&$ Baykal 1998). Using the archival RXTE observations, we measured three new pulse periods covering a time span of over 2 years in addition to the previous four pulse period measurements. With $\\sim 10^{3}-10^{4}$ sec observations separated by intervals of the order of a month we have been able to connect the pulses in phase and to construct the timing solution extending over a year. The residuals of pulse arrival times yielded a very low noise strength. Our findings imply that the source has a very stable spin down rate even over short time intervals, in contrast to the noise seen in other HMXRBs. ",
        "conclusions": ""
    },
    "0011/astro-ph0011318_arXiv.txt": {
        "abstract": "We present a method for calculating the infrared emission from a population of dust grains heated by starlight, including very small grains for which stochastic heating by starlight photons results in high temperature transients. Because state-to-state transition rates are generally unavailable for complex molecules, we consider model PAH, graphitic, and silicate grains with realistic vibrational mode spectra and realistic radiative properties. The vibrational density of states is used in a statistical-mechanical description of the emission process. Unlike previous treatments, our approach fully incorporates multiphoton heating effects, important for large grains or strong radiation fields.  We discuss how the ``temperature'' of the grain is related to its vibrational energy. By comparing with an ``exact'' statistical calculation of the emission process, we determine the conditions under which the ``thermal'' and the ``continuous cooling'' approximations can be used to calculate the emission spectrum.  We present results for the infrared emission spectra of PAH grains of various sizes heated by starlight. We show how the relative strengths of the 6.2, 7.7, and 11.3$\\micron$ features depend on grain size, starlight spectrum and intensity, and grain charging conditions. We show results for grains in the ``cold neutral medium'', ``warm ionized medium'', and representative conditions in photodissociation regions. Our model results are compared to observed ratios of emission features for the Milky Way and other galaxies, and for the M17 and NGC 7023 photodissociation regions. ",
        "introduction": "Very small grains -- consisting of tens to hundreds of atoms -- play a major role in the infrared emission from the interstellar medium. These grains are small enough that the time-averaged vibrational energy $\\langle E\\rangle$ is smaller than or comparable to the energy of the starlight photons which heat the grains. Stochastic heating by absorption of starlight therefore results in transient ``temperature spikes'', during which much of the energy deposited by the starlight photon is reradiated in the infrared.  The idea of transient heating of very small grains was first introduced by Greenberg (1968). Its importance was not apparent until the detection of the near infrared emission of reflection nebulae (Sellgren et al. 1983) and detection by the {\\it Infrared Astronomical Satellite (IRAS)} of 12 and 25$\\mu$m Galactic emission which was far in excess of the emission expected for interstellar dust at $T\\approx20\\K$ (Boulanger \\& P\\'erault 1988). Subsequent measurements by the {\\it Diffuse Infrared Background Experiment (DIRBE)} instrument on the {\\it Cosmic Background Explorer (COBE)} satellite confirmed this and detected additional broadband emission at 3.5 and 4.9$\\micron$ (Arendt et al.\\ 1998). More recently, spectrometers aboard the {\\it Infrared Telescope in Space (IRTS)} (Onaka et al. 1996; Tanaka et al. 1996) and the {\\it Infrared Space Observatory (ISO)} (Mattila et al. 1996) have shown that the diffuse interstellar medium radiates strongly in emission features at 3.3, 6.2, 7.7, 8.6, and 11.3$\\micron$. These emission features, previously observed from a small number of bright reflection nebulae, planetary nebulae, and HII regions, have been tentatively identified as emission from polycyclic aromatic hydrocarbons, or PAHs (L\\'eger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985).  To understand these observations we need to be able to calculate the infrared emission expected for a given model of interstellar dust. The question we address here is: what approximations can be used to carry out accurate calculations of thermal emission from a population of very small grains exposed to stochastic heating by starlight?  There have been a number of previous studies of the discrete heating of very small grains (Duley 1973; Greenberg \\& Hong 1974; Gail \\& Seldmayr 1975; Greenberg 1976; Purcell 1976; Drapatz \\& Michel 1977; Andriesse 1978; Aannestad \\& Kenyon 1979; L\\'eger \\& Puget 1984; Draine \\& Anderson 1985; Puget et al. 1985; D\\'esert et al. 1986; Dwek 1986; Guhathakurta \\& Draine 1989; Aannestad 1989; Siebenmorgen et al. 1992; Manske \\& Henning 1998; Duley \\& Poole 1998). Most studies have made three approximations: (1) the ``thermal'' approximation for excitation; (2) use of ``bulk'' specific heats; and (3) the ``continuous cooling'' approximation. While accurate for large grains, these approximations could potentially introduce significant inaccuracies for very small grains.  Barker \\& Cherchneff (1989) and d'Hendecourt et al. (1989) found that the thermal approximation was valid for computing thermal emission spectra. Allamandola, Tielens, \\& Barker (1989), Schutte, Tielens, \\& Allamandola (1993), and Cook \\& Saykally (1998) compared calculations of infrared emission from PAHs using the thermal approximation with the results of a more detailed statistical treatment; they found that the thermal approximation provided a good approximation for the emission from low-frequency modes, although it overestimated the emission from high frequency modes such as the $3.3\\micron$ C-H stretching mode. All of these studies used the ``continuous cooling'' approximation to discuss the overall emission from a PAH following single-photon heating.  The validity of the ``continuous cooling'' approximation has been questioned by Siebenmorgen et al. (1992), Manske \\& Henning (1998), and Duley \\& Poole (1998). Siebenmorgen et al.\\ (1992) and Manske \\& Henning (1998) attempted to treat the cooling process as discrete events, but did not correctly evaluate the transition probabilities, as we show below in \\S\\ref{sec:thermal-discrete}.  The present paper addresses the vibrational excitation of interstellar grains, both small and large, and the resulting infrared emission. If energy levels and transition probabilities were known, we could (at least in principle) solve for the statistical steady-state populations of the different energy levels of grains illuminated by a known radiation field. However, this level of detailed information is generally unavailable, for even the smallest and simplest PAH molecules. We therefore construct a ``model PAH'' grain, C$_n$H$_m$, with a realistic spectrum of energy levels. Following previous workers (e.g., Schutte et al.\\ 1993, Boulanger et al.\\ 1998) we adopt realistic absorption cross section $C_\\abs(\\lambda)$, depending on the number of C atoms, the H/C ratio, and whether the PAH is neutral or ionized. We construct a similar model for silicate grains, with appropriate mode spectrum and absorption cross section. We divide the energy level spectrum into a manageable number of ``bins'', estimate the ``bin-to-bin'' transition probabilities, and can then solve for the steady-state probability $P_j$ of occupying energy bin $j$.  Our analysis allows for discrete quantum states at low vibrational energies, and what is effectively a continuum of vibrational levels at large energies. As we show, the only simplifying assumptions that need to be made are: (1) {\\it the absorption cross section $C_\\abs(\\lambda)$ is independent of the degree of vibrational excitation}, and (2) {\\it internal vibrational relaxation rapidly redistributes internal energy among the vibrational degrees of freedom of the grain.} While not exact, these are both expected to be good approximations. Once $C_\\abs(\\lambda)$ is adopted and the spectrum of vibrational ``normal modes'' (and from this the vibrational density of states) is specified, the physics of absorption and emission of radiation by the grain is essentially fully determined. With these assumptions, we can solve for the energy distribution function $P(E)$ for a given grain in an arbitrary radiation field, with full inclusion of ``multiphoton heating'' effects which become important for large grains or intense radiation fields. Because we solve for $P(E)$ including all radiative transitions, we obtain the ``exact'' solution against which we can then test solutions obtained using the thermal approximation and the continous cooling approximation.  In \\S\\ref{sec:vibrational_mode_spectrum} we estimate the spectrum of vibrational modes for small PAH and silicate grains, and obtain the vibrational density of states for such grains. The transition probabilities depend on the photon absorption cross sections, and in \\S\\ref{sec:C_abs} we describe the cross sections adopted for PAH grains.  We formulate the ``exact'' statistical treatment in \\S\\ref{sec:statistical_treatment}. In \\S\\ref{sec:E(T)} we discuss how the ``temperature'' of a grain is related to its vibrational energy content, and in \\S\\ref{sec:thermal_treatment} we show how the ``thermal'' approximation is related to the exact treatment. The radiative cooling time for vibrationally-excited grains is estimated as a function of $E$ and grain size in \\S\\ref{sec:cooling_time}.  In \\S\\ref{sec:solution_method} we discuss the method used to solve for the steady-state energy distribution. In \\S\\ref{sec:IR_emission} we show how the infrared emission spectrum can be calculated from the energy distribution, both for the exact statistical treatment and using the thermal approximation. We show some sample solutions in \\S\\ref{sec:results}, where we assess the accuracy of the ``thermal'' approximation. We find the thermal approximation to be quite accurate for calculation of the overall emission spectrum from a population of very small grains. Furthermore, we find that the ``continuous cooling'' approximation is itself quite accurate, and can therefore be used to greatly reduce the amount of computation required to find the temperature distribution functions.  The above studies therefore provide a methodology for computing the infrared emission from dust grains irradiated by starlight. The method is accurate, and can be used for any specified mixture of grain sizes and compositions; we present results for both carbonaceous and silicate materials. We discuss our results in \\S\\ref{sec:discussion} and summarize the main points in \\S\\ref{sec:summary}.  The paper is largely concerned with the physics of infrared emission from stochastically-heated grains, and practical methods for calculating the emission spectrum.  Readers who are interested only in the resulting emission spectrum may wish to proceed directly to \\S\\ref{sec:results}. In a following paper (Li \\& Draine 2001a) we apply this methodology to construct models to account for the observed infrared emission from the diffuse interstellar medium. ",
        "conclusions": "}  \\subsection{Band Ratios}  From Figs.\\ \\ref{fig:IR_stat_exact} and \\ref{fig:IR_discrete_continuous} it is evident that the relative strength of the different PAH emission bands (3.3, 6.2, 7.7, 8.6, 11.3$\\micron$) is a strong function of the PAH size: small PAHs radiate strongly at 6.2 and 7.7$\\micron$, while larger PAHs emit most of their power at increasingly long wavelengths.  The 6.2, 7.7, and 11.3$\\micron$ features are prominent in many astronomical spectra. Let $P(\\lambda_0)$ be the power radiated in the feature with central wavelength $\\lambda_0$: for a Drude profile, $P(\\lambda_0)=(\\pi/2)(\\Delta P_\\lambda)_0 ({\\rm FWHM})_{\\lambda_0}$, where $(\\Delta P_\\lambda)_0$ is the peak contribution of the feature to the emission, and $({\\rm FWHM})_{\\lambda_0}$ is the full-width at half-maximum. In Figure \\ref{fig:band_ratios_0} we show how the relative strengths of these emission bands vary depending upon the size and charge state of the PAHs, and upon the starlight intensity $u_E$, characterized by $\\chi$ or $G_0$.\\footnote{ $\\chi$ is the intensity at 1000$\\Angstrom$ relative to the estimate of Habing (1968); $G_0$ is the ratio of the 6--13.6~eV energy density to the estimate of Habing (1968). The MMP spectrum has $G_0=0.923\\chi$; a $3\\times10^4\\K$ blackbody has $G_0=0.602\\chi$. } Neutral PAHs have large ratios of $P(11.3\\micron)/P(7.7\\micron)$, while PAH ions have much smaller values. For both neutrals and ions, there is a regular progression of $P(6.2\\micron)/P(7.7\\micron)$ to smaller values with increasing $N_\\C$. For a given PAH neutral or ion, the emission band ratios are essentially independent of $\\chi$ for $\\chi < 3\\times10^4(10^3/N_\\C)^{5/3}$ [see eq.\\ (\\ref{eq:N_sph})]; for larger values of $\\chi$, multiphoton heating effects begin to be evident. For $\\chi=10^6$, multiphoton effects are evident even for $N_\\C=10^2$ [consistent with eq.\\ (\\ref{eq:N_sph})].  \\begin{figure*}[ht] \\begin{center} \\hspace*{-15cm} \\begin{rotate}{-90} \\epsfig{file= f16.cps, width=11cm} \\end{rotate}\\vspace*{10cm} \\end{center} \\figcaption{\\footnotesize \\label{fig:band_ratios_0} Relative strengths of three PAH emission features (6.2 and 7.7$\\mu$m C-C stretching modes, and 11.3$\\mu$m C-H out-of-plane bending mode) for neutral and ionized PAHs, labelled by the number of carbon atoms. Results are shown for grains illuminated by the MMP spectrum with $\\chi=1.23$ and $123$, and for a $3\\times10^4\\K$ blackbody spectrum (cutoff at 912$\\Angstrom$) with $\\chi=10^2$, $10^4$, $10^5$, and $10^6$. Also shown (stars) are observed band ratios for the diffuse ISM, the reflection nebulae NGC 7023 and vdB 133, the M17 PDR, the $\\rho$ Oph molecular cloud; the starburst galaxies M82 and NGC 253 (St\\\"urm et al.\\ 2000); the Seyfert 2 galaxy in Circinus; the average of 28 normal galaxies; and a quiescent molecular cloud in the SMC. (See text for citations.) The dotted line connecting the neutral and ionized points for $N_\\C=16$ (in a PDR) illustrates the variation in the relative strengths of the 6.2, 7.7, and 11.3$\\micron$ features in a PDR as the PAH ion fraction varies from 0 to 1. } \\end{figure*} \\begin{figure*}[ht] \\begin{center} \\hspace*{-15cm} \\begin{rotate}{-90}\\epsfig{file= f17.cps, width=11cm} \\end{rotate}\\vspace*{10cm} \\end{center} \\figcaption{\\footnotesize \\label{fig:band_ratios_1} Relative strengths of three PAH emission features (6.2 and 7.7$\\mu$m C-C stretching modes, and 11.3$\\mu$m C-H out-of-plane bending mode); PAHs are labelled by the number of carbon atoms. Results are shown for 3 charging conditions: CNM, WIM, and PDR (see text). Circles: grains in CNM charging conditions, for $G_0=1.14$ and 114; diamonds: grains in WIM charging conditions, for $G_0=1.14$ and 114; squares: grains in PDR conditions, for $\\chi=10^2$, $10^4$, $10^5$, and $10^6$. Stars represent observations as in Fig.\\ \\ref{fig:band_ratios_0}. } \\end{figure*} \\begin{figure*}[ht] \\begin{center} \\epsfig{file=f18.cps,width=\\figwidthsmall} \\end{center} \\figcaption{\\footnotesize \\label{fig:band_ratios_2} Same as Fig.\\ \\ref{fig:band_ratios_0} but for a weighted sum of emission from grains in the cold neutral medium (CNM), grains in the warm neutral medium (WNM), and grains in the warm ionized medium (WIM), heated by the MMP estimate of the ISRF. Band ratios within the shaded region can be obtained by appropriate mixtures of PAHs with $20\\leq N_\\C\\leq 1000$. Observed band ratios for the diffuse ISM (Onaka et al. 1996; Mattila et al. 1996) and the average spectrum of 28 normal spiral galaxies (Helou et al. 2000) fall within this allowed region. } \\end{figure*}  In Figure \\ref{fig:band_ratios_0} we plot the observed ratios of $P(11.3\\micron)/P(7.7\\micron)$ and $P(6.2\\micron)/P(7.7\\micron)$ for (1) the diffuse ISM (Onaka et al.\\ 1996); (2) the average spectrum for a sample of 28 spiral galaxies (Helou et al.\\ 2000); (3) the reflection nebula at the surface of the $\\rho$ Oph molecular cloud (Boulanger et al.\\ 1996); (4) the reflection nebula vdB 133, illuminated by F5Iab and B7II stars (Uchida et al.\\ 1998); (5) the reflection nebula/photodissociation region (PDR) NGC 7023, illuminated by a B2.5V star (Cesarsky et al. 1996a); (6) the M17 PDR (Cesarsky et al.\\ 1996b); (7) the starburst galaxies M82 and NGC 253 (St\\\"urm et al.\\ 2000); (8) the Seyfert 2 galaxy in Circinus (St\\\"urm et al.\\ 2000); and (9) the quiescent molecular cloud SMC B1\\#1 in the Small Magellanic Cloud (Reach et al.\\ 2000). With the exception of SMC B1\\#1 (which we discuss below), the observed band ratios fall between the neutral and ionized PAH tracks in Figure \\ref{fig:band_ratios_0}, suggesting that the observed spectra can be reproduced by an appropriate mixture of PAH sizes, with an appropriate ionized fraction. The observed band ratios for the Milky Way ISM, $\\rho$ Oph, 28 normal galaxies, and the reflection nebulae vdB 133 and NGC 7023 could be reproduced with a relatively low ionized fraction, but larger ionized fractions are required for the M17 PDR, M82, NGC 253, and the Circinus galaxy.  \\subsection{SMC B1\\#1}  The quiescent molecular cloud SMC B1\\#1 observed by Reach et al.\\ (2000) falls outside the region bounded by the neutral and ionized tracks in Figure \\ref{fig:band_ratios_0}. Reach et al.\\ point out that SMC B1\\#1 has a 11.3/7.7 band ratio which differs significantly from that for Milky Way objects, including the $\\rho$ Oph reflection nebula (which one might expect SMC B1\\#1 to resemble). Reach et al.\\ suggest that the enhanced emission in the 11.3$\\micron$ C-H bending mode might be due to increased hydrogenation of the PAHs in the SMC, where the gas phase H/C ratio is a factor of 10 higher than in the Milky Way. Given that Milky Way PAHs with $\\gtsim$25 C atoms are expected to be essentially fully hydrogenated (Tielens et al.\\ 1987; Allamandola, Tielens, \\& Barker 1989; Allain, Leach, \\& Sedlmayr 1996), it is not clear that greater hydrogenation would be expected in the SMC. As discussed in Paper II, our model calculations assume that interstellar neutral and ionized PAHs have the 11.3/7.7$\\micron$ band strength ratio {\\it reduced} by a factor of 2 relative to PAHs which have been studied in the lab -- enhancement of the 7.7$\\micron$ feature appears to be necessary to reproduce the low 11.3/7.7$\\micron$ band ratios observed in objects such as NGC 253 or M82, and (with the possible exception of SMC B1\\#1) the reduced 11.3/7.7$\\micron$ band strength ratio appears to be consistent with astronomical observations.  While SMC B1\\#1 appears to fall ``outside'' the ``allowed'' region in Figures \\ref{fig:band_ratios_0}--\\ref{fig:band_ratios_2}, we show elsewhere (Li \\& Draine 2001b) that the present model can reproduce the observed line ratios by adding emission from a population of larger PAHs -- with $N_\\C\\approx 3000$ -- which radiate strongly in the 11--13$\\micron$ region, but produce little emission in the 6--8$\\micron$ region. These larger PAHs would radiate strongly at 15--50$\\micron$; SIRTF will be able to test this hypothesis.  \\subsection{PAH Ionization}  The fraction of PAHs which are charged will depend upon the electron density $n_e$, temperature $T$, and the starlight intensity. We consider three different charging environments: \\begin{itemize} \\item ``Cold neutral medium'' (CNM; $n_e=0.045\\cm^{-3}$, $T=100\\K$, and the MMP ISRF spectrum with $\\chi=1.23$, $G_0=1.14$) and with both radiation field and gas density increased by $10^2$; \\item ``Warm ionized medium'' (WIM; $n_e=0.1\\cm^{-3}$, $T=8000\\K$, and the MMP ISRF with $\\chi=1.23$), and with both radiation field and gas density increased by $10^2$; \\item Photodissociation region conditions (PDR; $n_e/n_\\rmH=3\\times10^{-4}$, $T=10^3\\K$, and a $3\\times10^4\\K$ blackbody cut off at 912$\\Angstrom$ with $\\chi/n_\\rmH=0.1\\cm^3$) for $\\chi=10^2$, $10^4$, $10^5$, and $10^6$. While PDRs come in a range of conditions, recent observations and modeling suggest that at least some PDRs include an extended region with $T\\approx10^3\\K$, and with $\\chi/n_\\rmH\\approx0.1\\cm^3$, as in a recent model for the NGC 2023 PDR (Draine \\& Bertoldi 2000). \\end{itemize} For each set of conditions, we consider PAHs with $N_\\C\\geq16$ C atoms, and include emission from both neutral and ionized forms, with the neutral fraction calculated assuming a balance between collisional charging by electrons and photodetachment and photoionization (Weingartner \\& Draine 2001b). The charging of small grains is largely a function of the quantity $\\chi\\sqrt{T}/n_e$ (Bakes \\& Tielens 1994; Weingartner \\& Draine 2001b). For the CNM, WIM, and PDR conditions considered here, $\\chi\\sqrt{T}/n_e$ = 270, 1100, and $10^4$ $\\cm^3\\K^{1/2}$, so we expect grains of a given size to be increasingly positively charged as we move from CNM to WIM to PDR.  The ``diffuse ISM'' and ``28 Galaxies'' points fall between the CNM and WIM tracks in Figure \\ref{fig:band_ratios_0}, implying that one could reproduce the overall emission from normal spiral galaxies by a weighted sum of interstellar conditions ranging from CNM to WIM, with a mixture of grain sizes such that the 6.2 -- 11.3$\\micron$ emission is dominated by grains with $20\\ltsim N_\\C \\ltsim 500$.  The observed flux ratios for NGC 7023 and M17 fall between the ``WIM'' and ``PDR'' tracks in Figure \\ref{fig:band_ratios_1}. The observed flux ratios for the starburst galaxies M82 and NGC 253, as well as the Circinus Seyfert 2 galaxy, fall close to the ``PDR'' tracks, consistent with the expectation that the spectra of these galaxies includes a substantial contribution from PDRs.  It is important to note that the ``PDR'' models in Figure \\ref{fig:band_ratios_1} have charging conditions such that a substantial fraction of the small grains are still neutral: 22\\% of the $N_\\C=100$ grains are neutral. Larger charged fractions -- and smaller values of $P(11.3\\micron)/P(7.7\\micron)$ -- can be achieved by increasing the value of $\\chi/n_\\rmH$. For example, if $\\chi/n_\\rmH$ is increased by a factor 2, only 10\\% of the $N_\\C=100$ grains are neutral.  In Figure \\ref{fig:band_ratios_2} we show the emission expected from grains of various sizes, with 43\\% in the CNM, 14\\% in the WIM, and 43\\% in the WNM ($n_e=0.03\\cm^{-3}$, $T=6000\\K$, MMP ISRF, $\\chi\\sqrt{T}/n_e=3200$). The shaded region shows the flux ratios which could be obtained by grain mixtures containing only grains $20<N_\\C<1350$ (from Figure \\ref{fig:IR_discrete_continuous} it is apparent that grains with $N_\\C>10^3$ are not likely to contribute significant emission at $\\lambda \\leq 11.3\\micron$). The observed flux ratios for the diffuse ISM in the Milky Way, and for the average of 28 spiral galaxies, both fall in the shaded region, suggesting that with an appropriate size distribution, one might be able to reproduce the observed emission. The question of size distributions will be addressed in a future paper (Li \\& Draine 2001a), where it will be shown that the observed infrared emission can be approximately reproduced by grain size distributions similar to those invoked to explain the observed microwave emission from interstellar dust (Draine \\& Lazarian 1998a,b).  } We have presented a method for calculating the infrared emission from dust grains, including very small PAH molecules or silicate clusters. The principal results of this paper are the following: \\begin{enumerate} \\item The vibrational mode spectrum and vibrational density of states of very small grains are discussed in \\S\\ref{sec:vibrational_mode_spectrum}, where we explicitly consider both polycyclic aromatic hydrocarbons and silicate grains. \\item We use the adopted vibrational mode spectrum to calculate the specific heat of carbonaceous and silicate grains; the results are in agreement with experimental specific heats for graphite and amorphous silicates (Figure \\ref{fig:spec_heats}). \\item Using estimates for the absorption cross sections of carbonaceous and silicate grains, we obtain (in \\S\\ref{sec:statistical_treatment}) radiative transition rates for emission from vibrationally-excited states of carbonaceous and silicate grains. We refer to this as the ``exact-statistical'' treatment. \\item We discuss (in \\S\\ref{sec:thermal_treatment}) how the ``thermal approximation'' can be used to estimate rates for spontaneous emission. With a judicious estimate for the ``temperature'', the thermal approximation provides transition rates which are close to those obtained from the ``exact-statistical'' treatment. Therefore level populations estimated using the thermal approximation are close to the actual level populations given by the ``exact-statistical'' treatment. \\item We discuss the ``continuous cooling'' approximation for modeling the deexcitation of the grain. We show that level populations and emission spectra computed using this approximation are sufficiently accurate for most astrophysical applications, even for grains containing as few as $\\sim30$ atoms. \\item Methods for numerical solution for the vibrational energy distribution are presented in \\S\\ref{sec:solution_method}, and the numerical advantages of the continuous cooling approximation are stressed. \\item Relative strengths of the 6.2, 7.7, and 11.3$\\micron$ features depend on the grain size, on the charging conditions, and on the starlight intensity.  Observed 6.2/7.7 and 11.3/7.7 band ratios for the diffuse ISM in the Milky Way, for the $\\rho$ Oph and vdB133 reflection nebulae, and for the 28 ``normal'' galaxy sample of Helou et al. (2000), appear consistent with grains in the diffuse ISM heated by the average starlight background. The band ratios observed for the starburst galaxies M82 and NGC 253, for the Seyfert 2 ``Circinus galaxy'', for the M17 PDR, and for the NGC 7023 PDR appear to require conditions close to our ``PDR'' conditions, with $\\chi/n_\\rmH\\approx 0.05 - 0.1\\cm^3$. \\item The unusual band strengths observed for the quiescent molecular cloud SMC B1\\#1 in the SMC (Reach et al.\\ 2000) cannot be reproduced by a mixture of small $N_\\C\\ltsim 10^3$ PAHs. Either the PAHs in this cloud differ from Milky Way PAHs in their properties, or have a different size distribution. \\end{enumerate} The present paper has concentrated on the statistical mechanics of the stochastic heating process, and on the resulting emission spectra as a function of grain size. In a separate paper (Li \\& Draine 2001a) we calculate the emission from grain models with realistic size distributions, and show that the observed infrared emission from interstellar dust can be reproduced by a grain model with $\\sim$15\\% of the interstellar carbon in PAH particles containing less than $10^3$ C atoms."
    },
    "0011/astro-ph0011410_arXiv.txt": {
        "abstract": "We present an analysis of the frequency shift and the even terms of the frequency splitting coefficients carried out using six years of LOWL data, starting in 1994. The temporal variations, and their dependences with the frequency and degree are addressed. The results are consistent with previous analysis. ",
        "introduction": "It is very well established that the variations of some of the oscillation mode parameters is a  signature of the solar activity cycle. For instance, the acoustic mode frequency is shifted positively of 0.4$\\mu$Hz from minimum to maximum of the solar cycle (Libbrecht and Woodard 1990, R\\'egulo  et al. 1994, Jim\\'enez-Reyes  et al. 1998). It has also been noted that, this variation depends on the frequency and on the degree of the mode as well. The even terms of the splitting coefficients present similar changes. Here, we present an analysis of the frequency shift using LOWL observations. ",
        "conclusions": "We have analyzed six years of LOWL data from 1994 to the end of 1999 in order to parameterize the variation of the frequency shift and the even $a$-coefficients.  In summary, the acoustic central frequency presents a variation which is very well correlated with the solar cycle. The integrated amplitude of this variation, peak-to-peak, is close to 0.4$\\mu$Hz. This variation is however a function of both the frequency and the degree, and these dependences can be well fitted by the inverse mode mass.  In addition, the even $a$-coefficients present a significant change with the solar cycle, very well correlated with the solar indices. Moreover, as for the frequency shift, these changes increase with the frequency. The results agree very well with earlier studies (Libbrecht and Woodard, 1990) carried out at the beginning of the last solar cycle.  LOWL instrument has been recently updated and a new instrument has been developed and installed at the Observatorio del Teide (Tenerife), leading to a new network called Experiment for Coordinated Helioseismic Observations (ECHO). ECHO is intended to continue observing for a complete solar cycle allowing us to better track the origin of the solar activity cycle."
    },
    "0011/astro-ph0011283_arXiv.txt": {
        "abstract": "We consider disk accretion onto a magnetic star under the assumption that the stellar magnetic field permeates the disk and that the magnetosphere that lies between the disk and the star is force free. Using simplified axisymmetric models (both semianalytic and numerical), we study the time evolution of the magnetic field configuration induced by the relative rotation between the disk and the star. We show that, if both the star and the magnetosphere can be approximated as being perfectly conducting, then the behavior of the twisted field lines depends on the magnitude of the surface conductivity of the disk. For any given relative azimuthal speed $\\Delta v_\\phi$ between the disk and the star (measured at the disk surface), there is a maximum surface conductivity $\\Sigma_{\\rm max} \\sim c^2/|\\Delta v_\\phi|$ such that, if the actual surface conductivity is smaller than $\\Sigma_{\\rm max}$, then a steady-state field configuration can be established, whereas for larger values no steady state is possible and the field lines inflate and open up when a critical twist angle (which for an initially dipolar field is $\\sim 2\\ {\\rm rad}$) is attained. We argue that thin astrophysical disks are likely to have surface conductivities that exceed the local $\\Sigma_{\\rm max}$ except in regions where $|\\delta v_\\phi|$ is particularly small (such as the immediate vicinity of the corotation radius in a Keplerian disk).  We find that, if the disk conductivity is high enough for the twist angle to continue growing until the field lines open, then, except under rather special circumstances, the radial magnetic field at the disk surface will grow to a level for which the radial migration of the field lines in the disk cannot be ignored. A similar conclusion was reached by Bardou \\& Heyvaerts (1996). We demonstrate, however, that the radial diffusion in the disk is much slower than the field-line expansion in the magnetosphere, which suggests that, contrary to the claim by Bardou \\& Heyvaerts, the twisting process does {\\em not} result in a rapid expulsion of the field lines from the disk.  We calculate the magnetospheric density redistribution brought about by the field-line expansion and show that inertial effects in the magnetosphere will slow the expansion (and invalidate the assumption of a force-free field) before the critical twist angle is reached. We argue, however, that these effects merely slow down the opening process but do not inhibit it altogether. We find that the field opening drastically reduces the density near the apex of the expanding field lines (which typically elongate in a direction making an angle $\\ga 60^{\\circ}$ to the rotation axis) while also creating a pronounced density enhancement near the rotation axis. The former effect is conducive to the triggering of microinstabilities (e.g., ion-acoustic) in the $j_\\phi$ current layer that forms along the direction of elongation, whereas the latter is evidently related to a mechanism for the formation of an axially outflowing condensation that was previously identified in axisymmetric numerical simulations of such systems.  We examine the possibility that the expanding field lines reconnect across the current layer before they open up, so that the twisting leads to a quasi-periodic process of field-line expansion and reconnection and not just to a one-time opening event. We tentatively conclude that hyperresistivity associated with tearing-mode turbulence could lead to efficient reconnection, but we speculate that even faster reconnection might be brought about by 3-D kinking of the twisted field lines that gave rise to thin current sheets. This question could be further pursued through MHD stability calculations and 3-D numerical simulations. ",
        "introduction": "\\label{sec-intro}  The process of disk accretion onto a magnetic star is thought to have important consequences for the spin evolution and observational characteristics of neutron stars, white dwarfs, and young stellar objects (YSOs). In particular, if the stellar magnetic field lines penetrate the disk, then they may transmit torques between the disk and the star. Such an interaction has been invoked to explain the spin-up and spin-down episodes observed in X-ray pulsars (e.g., Ghosh \\& Lamb 1978, 1979a, 1979b, hereafter GL; Wang 1987; Lovelace, Romanova, \\& Bisnovatyi-Kogan 1995, hereafter LRBK; Yi, Wheeler, \\& Vishniac 1997) as well as the observed rotation-period distribution in YSOs (e.g., K\\\"onigl 1991; Edwards et al. 1993; Bouvier et al. 1993; Collier Cameron \\& Campbell 1993; Yi 1994, 1995; Ghosh 1995; Collier Cameron, Campbell, \\& Quaintrell 1995; Herbst et al. 2000). Furthermore, if the field is strong enough, then it may truncate the disk before it reaches the stellar surface and channel the inflowing matter to high stellar latitudes, where the (by now nearly free-falling) gas is stopped and thermalized in accretion shocks. This is the accepted explanation for X-ray pulsars (accreting magnetic neutron stars; e.g., Lamb 1989) and DM Herculis stars (accreting magnetic white dwarfs; e.g., Patterson 1994), and it has also been proposed as the origin of the optical/UV ``hot spots'' in classical T Tauri stars (interpreted as accreting magnetic YSOs; e.g., Bertout, Basri, \\& Bouvier 1988; K\\\"onigl 1991; Edwards et al. 1994; Hartmann, Hewett, \\& Calvet 1994; Lamzin 1995; Bertout et al. 1996; Johns-Krull \\& Basri 1997; Johns-Krull \\& Hatzes 1997; Martin 1997; Muzerolle, Hartmann, \\& Calvet 1998).  However, because of the complexity and diversity of the physical processes that are involved, the theory of the interaction between a magnetic star and a surrounding accretion disk remains incomplete. One key question is whether a steady-state description, which has been adopted in many of the models constructed to date, is appropriate. In a Keplerian disk around a star of mass $M$ that rotates with angular velocity $\\Omega_*$, the gas interior to the {\\em corotation radius} $r_{\\rm co} = (GM/\\Omega_*^2)^{1/3}$ rotates faster than the star, whereas the matter at larger radii rotates more slowly. If the star, the disk, and the magnetosphere that occupies the space between them can be treated as perfect conductors, then stellar magnetic field lines that thread the disk at any radius other than $r_{\\rm co}$ will undergo secular twisting. One possibility for establishing a steady state is for the twisting of the field lines to be countered by magnetic diffusivity in the disk (e.g., GL). However, in many cases a realistic estimate of the disk diffusivity leads to values that are much too small to justify a steady-state description of a system with a dipolar field configuration except in the immediate vicinity of $r_{\\rm co}$ (see \\S~\\ref{subsec-resistive-radial} below).  The behavior of twisted magnetic field lines that are anchored in a well-conducting medium has originally been considered in the context of the solar corona by means of semianalytic (e.g., Aly 1985; Low 1986; Wolfson 1994; Aly 1995) and numerical (e.g., Miki\\'c \\& Linker 1994; Amari et al. 1996a,b) techniques. However, direct insight into the field evolution in rotating accretion disks has by now been gained also from explicit investigations of magnetically linked star--disk systems, in which both semianalytic [e.g., van Ballegooijen 1994 (hereafter VB); Lynden-Bell \\& Boily 1994; LRBK; Bardou \\& Heyvaerts 1996 (hereafter BH)] and numerical (e.g., Hayashi, Shibata, \\& Matsumoto 1996; Goodson, Winglee, \\& B\\\"ohm 1997; Miller \\& Stone 1997; Goodson, B\\\"ohm, \\& Winglee 1999; Goodson \\& Winglee 1999) approaches have again been employed. These studies have indicated that the applied twist leads to a strong expansion of the field lines away from the star. Initially, the evolution is quasi-static, with the azimuthal field component building up while the poloidal field structure changes relatively little, but past a certain point the expansion accelerates rapidly and in a finite time (corresponding to a total twist $\\sim \\pi$) approaches a singular state characterized by the opening of at least a fraction of the field lines.  The evolution of the magnetic field lines after the singular state is reached is another important question on which there has been no unanimity in the literature. According to one school of thought (e.g., LRBK), the opening effectively severs the magnetic link between the disk and the star once and for all. The opposite view is that the star--disk link is reestablished through rapid field-line reconnection in the magnetosphere, and that the field then continues to undergo successive episodes of twisting, expansion, and reconnection (e.g., Aly \\& Kuijpers 1990, Goodson et at 1997, 1999). The latter picture is consistent with the results of numerical simulations of twisted coronal arcades, in which the addition of finite resistivity was found to result in reconnection across the current sheet that forms in that case (e.g., Miki\\'c \\& Linker 1994; Amari et al. 1996a). It has, furthermore, been argued that, in a real magnetosphere, the current concentration will itself facilitate reconnection through the initiation of plasma microinstabilities that give rise to an anomalous resistivity (e.g., Lamb, Hamilton, \\& Miller 1993). To resolve this question in a quantitative manner, it is necessary to calculate the evolution of the magnetospheric current and particle densities as the field lines are being twisted in order to determine whether an instability is likely to develop before the field expands to such a degree that it becomes effectively open.  When the expanding, twisted field lines approach the singular state, they typically become sharply bent at the disk surface, giving rise to a strong radial stress that tends to push the field outward through the disk for reasonable values of the disk diffusivity. Previous treatments of the problem have either ignored this issue altogether (e.g., LRBK), concluded that it would lead to the effective expulsion of the field lines from the disk (BH), or else considered the intermediate case in which the radial field excursions average to zero (VB). The resolution of this issue depends in part on whether reconnection in the magnetosphere terminates the field-line expansion phase before the radial stress at the disk surface becomes very large, and it is thus directly related to the field-opening question.  Even if the field configuration is not steady, the duration of the successive expansion/reconnection cycles is sufficiently short [of order the dynamical (rotation) time], that one could consider averaging the relevant quantities over the cycle period (cf. VB) and using them in modeling effectively steady-state star--disk systems (e.g., GL; Zylstra 1988; Daumerie 1996). However, the time-dependent nature of the magnetic interaction could be crucial to the origin of the observed accretion and outflow properties of disk-fed magnetic stars (e.g., Hartmann 1997; Goodson \\& Winglee 1999), and it may also help resolve some of the difficulties identified in steady-state magnetic accretion models (e.g., Safier 1998).  In this paper we address various aspects of the basic issues listed above using a simple, largely semianalytic, approach. Our modeling framework, based on evolving a series of self-similar, force-free equilibrium magnetospheric field configurations, is described in \\S~\\ref{sec-model}. In \\S~\\ref{sec-resistive} we clarify the condition for steady-state evolution in a disk with a given distribution of magnetic flux and diffusivity, and we then consider also the radial migration of flux across the disk. In \\S~\\ref{sec-magnetosphere} we examine the opening of twisted magnetic field lines using a uniformly rotating, self-similar disk model, and we investigate the role of reconnection and inertial effects in the magnetosphere above the disk in limiting this process. In \\S~\\ref{sec-keplerian} we generalize our analysis of the opening of magnetic field lines by using a non--self-similar model that incorporates a differentially rotating disk. In \\S~\\ref {sec-discussion} we discuss our results and relate them to recent numerical simulations. We summarize in \\S~\\ref{sec-summary}. ",
        "conclusions": "\\label{sec-discussion}  In this section we consider the implications of our work to the various modeling issues discussed in \\S~\\ref{sec-intro}. We start with one of our key results, namely, the finding (\\S~\\ref{sec-resistive}) that real astrophysical systems are unlikely to reach an exact steady state unless they have significant {\\em negative} radial field components at the disk surface in their untwisted (potential) state (corresponding to an $n\\ll 1$ magnetic field configuration in the self-similar representation). We now argue that such low-effective-$n$ configurations are not likely to be realized in axisymmetric, magnetically linked star--disk systems.  Although models in which the stellar magnetic field is pulled into the disk have been considered in the literature (e.g., Ostriker \\& Shu 1995), and $B_{{\\rm d},r}$ is $<0$ even for a strictly dipolar field if the disk has a finite height, it is difficult to imagine that strong negative radial field components would naturally develop in the course of disk accretion onto a magnetized star. In fact, simple global disk models (e.g., Spruit \\& Taam 1990) indicate that, so long as the disk is at least partially conductive, it will tend to exclude the (dipolar) stellar magnetic field and cause some of the magnetic flux to remain trapped inside its inner radius, thereby giving rise to a {\\em positive} radial field component at the disk surface. Furthermore, there has been no indication from any of the numerical simulations carried out to date that the magnetic field in a disk interacting with a central dipole can evolve into an $n\\ll 1$ configuration. Although the published simulations have all been strictly 2-D, such systems are even less likely in 3-D. This is because the steady-state values of $|B_{{\\rm d},\\phi}/B_{{\\rm d},z}|$ are expected to be well in excess of 1 for $n \\ll 1$ (see Table 1), implying that the field configuration might be susceptible to a disruptive, nonaxisymmetric instability (e.g., Kadomtsev 1966).  It is conceivable, however, that negative values of $B_{{\\rm d},r}$ would occur in the case of a {\\em nonaxisymmetric} field configuration. Nonaxisymmetry could be introduced by the misalignment of a stellar dipolar field axis with the disk axis, which is, in fact, the common interpretation of many of the observed periodicities in accreting magnetic stars (e.g., Lamb 1989; Patterson 1994; Bertout et al. 1988). Alternatively, it could arise from the presence of isolated stellar magnetic loops that thread the disk (e.g., Bertout et al. 1996; Safier 1998). The latter possibility is consistent with the inferred (by radio observations) existence of extended ($\\ga 10 \\, R_*$), organized magnetic structures in certain YSOs (e.g., Andr\\'e  et al. 1992). The axisymmetric self-similar model with $n\\, < \\, 1$ could provide a qualitative representation of the field evolution in this case, but some important aspects, such as the possibility that only a fraction of the field lines open first (e.g., Amari et al. 1996b), or the development of kink and/or buckling instabilities (e.g., Parker 1979), might be captured only in a fully 3-D calculation. In any case, although such systems might attain a steady time-averaged configuration, they could never be in a strictly steady state.  Another new result of our work has been the demonstration (\\S~\\ref{subsec-resistive-radial}) that, even in disks with a comparatively low diffusivity, in which the field lines are twisted until they open up, one cannot consistently neglect the {\\em radial} migration of the magnetic footpoints in the disk as the critical twist angle $\\Delta\\Phi_{\\rm c}$ is approached. We have argued that it is unlikely that a state with $B_{{\\rm d},r} = 0$ (even just in a time-averaged sense) can be attained unless, again, the magnetic field configuration corresponds to that of an $n\\ll 1$ self-similar model. Under these circumstances, a disk that undergoes cycles of successive opening and closing (by reconnection) of twisted field lines (and, to a lesser extent, even a disk in which a steady-state twist is maintained) should manifest a secular evolution of the radial flux distribution. We have explicitly shown, however, that the rate of outward flux diffusion in the disk is significantly smaller than the rate of expansion of the twisted field lines in the overlying magnetosphere. This provides a strong argument {\\em against} the suggestion made by BH that the field lines in low-diffusivity disks remain perpendicular to the disk surface at all times. BH's conclusion that the twisting process will tend to expel to the outer edge of the disk almost all the stellar field lines that initially threaded it can be questioned on similar grounds. Although an outward flux migration might take place, we consider it unlikely that it would lead to a flux rearrangement in the disk that is as drastic as envisioned by BH. In fact, in certain cases a steady time-averaged flux distribution may be established. This is evidently the case in the time-dependent model presented by Goodson et al. (1999) and Goodson \\& Winglee (1999), in which inward radial mass motions that counter the outward radial flux diffusion take place in the disk.  We have also investigated (\\S~\\ref{sec-magnetosphere}) the density redistribution in the magnetosphere brought about by the motion of the twisted field lines. We showed that, as the critical twist angle is approached, the density in the apex region of the elongated flux ropes drops precipitously, with some of the depleted mass being pushed toward the rotation axis, where it becomes concentrated in a narrow angular sector about $\\theta=0$. A similar axial density enhancement has been found in the numerical simulations of Goodson et al. (1997, 1999; see also Hayashi et al. 1996), where it is directly associated with the formation of an overdense, ``knotty'' jet. As already noted by Goodson et al. (1999), this mechanism of producing axially outflowing condensations requires a good alignment between the rotation and magnetic dipole axes. Although our sequence-of-equilibria representation cannot capture the complex dynamical evolution of the simulated jets, it is gratifying that this simple model nonetheless allows a quantitative description of the basic process that generates their density structure.\\footnote {The simulations by Hayashi et al. (1996) and Goodson et al. (1997, 1999) also show an equatorial density enhancement induced by mass being pushed away from the apex regions of the expanding flux loops toward lower latitudes. This feature of the density evolution is not reproduced by our calculations on account of the fixed-density boundary condition that we imposed at the disk surface (see~eq.~[4.20]).}  Although inertial effects slow down the field-line expansion, one can expect that the twisted field lines will still {\\em effectively open} in a finite time. By this we mean that for any region of a given size threaded by a finite amount of flux, the apexes of the majority of the field lines will leave the region under condsideration after being twisted by a finite angle. This is indeed the behavior found in numerical simulations. For example, in the model of Goodson et al. (1999), even though inertial effects are enhanced by a relatively high initial mass loading of the magnetosphere and by subsequent episodes of post-reconnective mass ejection from the disk, the field lines typically open in less than one differential star--disk rotation period (as stated in Goodson \\& Winglee 1999), which is to be compared with the time ($\\la 1/2$ period) predicted by our model in the absence of inertial effects (see \\S~\\ref{sec-model} and \\S~\\ref{sec-keplerian}). The basic question is then whether the field lines eventually reconnect, resulting in a repetitive cycle of inflation and reconnection, or whether the opening process is essentially a one-time affair.  The 2-D numerical simulations carried out to date do not provide a clear answer to this question. Although reconnection was found to take place in the work of Hayashi et al. (1996), this was probably a consequence of their adoption of an unrealistically low threshold value of $j/\\rho$ for the onset of anomalous resistivity and of a rather strong functional dependence [$\\eta \\propto (j/\\rho)^2$] above the threshold. Numerical diffusivity was evidently also responsible for the reconnection events observed in the simulations by Romanova et al. (1998) of twisted, {\\em disk-anchored} magnetic loops. In the case of the numerical model of Goodson et al. (1999), the nominal value of the numerical resistivity-based Lundquist number in the innermost zone was $\\sim 100$ at the start of their run. As noted in \\S~\\ref {subsec-reconnection}, our self-similar model results indicate that a value in this range is compatible with current-driven Sweet-Parker reconnection. However, the reconnection that actually took place in this simulation was triggered by inward mass motions at the inner radius of the disk that strongly pinched the magnetosphere (see also Miller \\& Stone 1997).  Our analysis of the 2-D self-similar model has led us to the conclusion that ion-acoustic microinstability will be triggered too late to result in significant reconnection before the field lines effectively open. The crux of the problem is that the current layer that forms in this geometry along the field-line elongation direction is comparatively wide (aspect ratio $\\sim 10$ for an initially dipolar field), so the effective Lundquist number that is required for efficient reconnection is rather low and is attained only relatively late in the field-line opening process. Tearing-mode turbulence could, in principle, lead to reconnection on the required scale if the tearing fluctuation level is sufficiently high. As discussed by Strauss (1988), such turbulence can occur in 3-D, giving rise to hyperresistivity that may result in fast reconnection. We can, in fact, argue more generally that, in contrast with the idealized 2-D configurations that we employed in our model, real 3-D systems may exhibit nonaxisymmetric behavior that could significantly broaden the range of possible scenarios. For example, it seems quite probable that, as the field lines are twisted, the 2-D equilibrium will become unstable to nonaxisymmetric MHD perturbations, such as the kink mode (a possibility already mentioned by LRBK). The detailed MHD stability analysis that is required for a closer examination of this suggestion is beyond the scope of this paper, but it seems likely that such instabilities could induce a transition to a new, nonaxisymmetric configuration that might have real (i.e., infinitesimally thin) current sheets. Such current sheets could form sites of rapid reconnection that would enable the system to relax to a less stressed state, and it is quite probable that the field-line expansion would then terminate much sooner than in the 2-D scenario considered above. This possibility could be checked by means of future 3-D numerical simulations.  We thus regard it as highly plausible that reconnection, leading to periodic (or quasi-periodic) evolution, would be the realistic outcome of the field-line twisting process.  This inherently nonsteady behavior enhances the attractiveness of this scenario as a modeling framework. In particular, in the case of YSOs, it could be relevant to the interpretation of such phenomena as X-ray flares, the formation of a ``knotty'' axial jet and of a less collimated and slower (but more massive) disk outflow, and the heating of the permitted line-emitting gas in the magnetosphere (e.g., Hayashi et al. 1996; Hirose et al. 1997; Miller \\& Stone 1997; Hartmann 1997; Shu et al. 1997; Goodson  et al. 1997, 1999; Goodson \\& Winglee 1999). Furthermore, since the successive openings and reconnections of the field lines have been inferred to modulate the gas inflow rate onto the star and the net torque exerted on the star by the accreting matter, they would also have a direct impact on the mass and spin evolution of the protostar. It may nevertheless be possible to construct stationary, global models of accretion disks in such systems by obtaining a suitable time average of the magnetic field configuration. This is, in fact, the approach originally undertaken by GL and subsequently elaborated upon by Zylstra (1988) and Daumerie (1996). In particular, Daumerie developed a model in which the flux distribution along the disk is calculated by approximating the magnetic field in the magnetosphere as being current-free and by using an effective conductivity (based on the value of $B_{{\\rm d},\\phi}/B_{{\\rm d},z}$ through a relation analogous to our eq. [3.5]) to evaluate the poloidal field components in the disk. We note, however, that inasmuch as the azimuthal currents (which determine the disk contribution to the poloidal magnetic field) flow entirely inside the disk, one should employ the actual disk conductivity in this calculation instead of the above ansatz, as the latter could significantly overestimate the correct value in regions where the effective surface conductivity exceeds the limit (3.6). It may also be possible to refine the treatment of the magnetospheric field by using the approach outlined in \\S~\\ref{sec-keplerian} to calculate the magnetic field configuration in the fully force-free formulation (i.e., without neglecting current flows in the magnetosphere). After identifying the time at which reconnection is most likely to occur, one could proceed to derive the time-averaged values of $B_{{\\rm d},r}(r)/B_{{\\rm d},z}(r)$ and $B_{{\\rm d},\\phi}(r)/B_{{\\rm d},z}(r)$ and use them as boundary conditions for the disk model.  Past numerical simulations have proved invaluable in clarifying some apparent difficulties and in pointing the way to further theoretical progress. For example, Safier (1998) has raised concerns about the ability of the magnetic field lines to thread the disk rather than be swept out in a thermally driven stellar wind. However, the numerical simulations (which typically do include thermal effects) have demonstrated that a field with an initially dipolar topology can maintain a configuration in which it continuously threads the disk even as some field lines open and drive an outflow, and that various time-dependent effects (such as field-line twisting and mass motions in and out of the disk plane) combine to increase its strength above naive, steady-state estimates. Nevertheless, more refined simulations (affording better control on the field diffusivity and incorporating 3-D effects) are needed to fully resolve this and other still-open issues.     Disk accretion onto a magnetic star occurs in a variety of astrophysical contexts, ranging from T Tauri stars to X-ray pulsars. The magnetohydrodynamic interaction between the stellar field and the accreting matter can have a strong effect on the disk structure, the transfer of mass and angular momentum between the disk and the star, and the production of bipolar outflows. In this paper we concentrated on a key element of this interaction --- the time evolution of the magnetic field lines that thread the disk resulting from the relative rotation between the disk and the star --- and studied it using simplified models of axisymmetric, force-free fields. In particular, we employed the semianalytic similarity solution first derived by van Ballegooijen (1994) and Lynden-Bell \\& Boily (1994) to construct a sequence of magnetospheric equilibria parametrized by the relative twist angle $\\Delta \\Phi$, but we tested the generality of the basic conclusions of this model (which strictly applies only to a uniform relative angular speed) by constructing numerical solutions of the Grad-Shafranov equation for systems with rotation laws that approximate a Keplerian disk. Our main results can be summarized as follows.  \\begin{itemize} \\item On the assumption that both the star and the magnetosphere can be approximated as being perfectly conducting, the behavior of the twisted field lines depends on the magnitude of the surface conductivity $\\Sigma$ of the disk. A steady-state configuration can be established only if, at a radius $r$ in the disk, $\\Sigma(r) < \\Sigma_{\\rm max}(r) \\sim c^2/r|\\Delta \\Omega(r)|$ (eq. [3.6]). Weakly ionized protostellar disks as well as disks characterized by a turbulent diffusivity typically do not satisfy this condition except very close to the corotation radius, implying that the field lines undergo secular inflation and will open up when a critical twist angle $\\Delta\\Phi_c$ ($ \\sim 2\\ {\\rm rad}$ for an initially dipolar field configuration) is reached.  \\item If inertial effects in the magnetosphere could be neglected, then the field lines would open in a fraction of the differential star--disk rotation period. However, inertial effects intervene, slowing down the expansion speed at the apex of the twisted field line to the local Alfv\\'en speed (which, for an initial dipolar field and a power-law density distribution, remains constant throughout the evolution in the self-similar model). This will delay the effective opening of the field lines, but will not by itself prevent it from occurring.  \\item If $\\Sigma(r) > \\Sigma_{\\rm max}(r)$ and the field lines undergo a strong expansion, then the radial magnetic stresses at the disk surface will cause the field lines to start migrating outward. A similar conclusion was reached by Bardou \\& Heyvaerts (1996), who went on to suggest that the field lines will be rapidly expelled from the disk.  We argued, however, that this is unlikely to happen in view of the fact that the radial diffusion in the disk is much slower than the field-line expansion in the magnetosphere. If the field lines expand and reconnect in a periodic manner then a time-averaged steady state in which the mean surface radial field is zero is in principle possible (van Ballegooijen 1994), but this situation will not arise if the initial field configuration is dipolar (although it might conceivably pertain to isolated stellar flux tubes that thread the disk).  \\item The opening of the field lines induces a redistribution of the magnetospheric density $\\rho$. In particular, it leads to a strong depletion of mass near the apex of the expanding field lines (which typically elongate in a direction making an angle $\\ga 60^{\\circ}$ to the rotation axis) and to a pronounced density enhancement near the rotation axis. The decrease in density along the direction of elongation (which also marks the location of a $j_\\phi$ current layer) is conducive to the triggering of microinstabilities (such as ion-acoustic), since the instability onset condition corresponds to a threshold value of $j/\\rho$. The density enhancement near the axis (which is largely a byproduct of the axisymmetry assumption) represents a possible mechanism for the formation of an axially outflowing condensation (generalizing to a ``knotty'' jet if the process is repetitive), as suggested previously (e.g., Goodson et al. 1999) on the basis of numerical simulations.  \\item The question of whether the evolution of the twisted field culminates in the effective opening of the field lines (e.g., Lovelace et al. 1995), or whether it is a repetitive process (e.g., Aly \\& Kuijpers 1990) of field-line expansion and reconnection (not necessarily back to exactly the same field configuration that existed before the most recent expansion, in which case it will not be strictly periodic), is still open. The reconnection events observed in the numerical simulations carried out to date are by and large due to numerical diffusivity and therefore do not provide conclusive answers. Based on an analysis of the 2-D self-similar model, we concluded that anomalous resistivity derived from the ion-acoustic/Buneman instability will be triggered too late for significant reconnection to take place across the comparatively wide (aspect ratio $\\sim 10$ for an initially dipolar field) current layer that forms in this case before the field lines effectively open. We suggested, however, that hyperresistivity associated with tearing-mode turbulence (Strauss 1988) might lead to efficient reconnection, although whether such turbulence could arise in the configuration under consideration is not clear. In addition, it is conceivable that even faster reconnection could take place as a result of 3-D kinking of the twisted field lines that would create genuine (i.e., infinitesimally thin) current sheets. This issue might be resolved by MHD stability calculations and 3-D numerical simulations.  \\end{itemize}  \\vskip 20pt  We thank Bob Rosner, Oded Regev, Aad van Ballegooijen, Fred Lamb, and B.C. Low for interesting and fruitful discussions. This research was supported in part by NASA grants NAG 5-3687 and NAG 5-1485 and by the ASCI/Alliances Center for Astrophysical Thermonuclear Flashes at the University of Chicago under DOE subcontract B341495.   \\appendix"
    },
    "0011/astro-ph0011556_arXiv.txt": {
        "abstract": "Deep long-slit spectra of the diffuse ionized gas halos of the edge-on spiral galaxies NGC 4302 and UGC 10288 are presented.  Emission lines are detected up to about $z=2$ kpc in NGC 4302, and to nearly $z=3$ kpc on the north side of UGC 10288.  For both galaxies, the line ratios [\\ion{N}{2}]/H$\\alpha$ and [\\ion{S}{2}]/H$\\alpha$ increase with $z$ in accordance with dilute photo-ionization models.  Runs of [\\ion{S}{2}]/[\\ion{N}{2}], and for UGC 10288, the run of [\\ion{O}{3}]/H$\\alpha$, however, are not explained by the models.  Scale height determinations of their DIG halos are generally lower than those of galaxies with more prominent extraplanar DIG features.  These data, along with previously presented data for NGC 5775 and NGC 891, are used to address the issue of how DIG halos are energized.  Composite photo-ionization/shock models are generally better at explaining runs of line ratios in these galaxies than photoionization models alone.  Models of line ratios in NGC 5775 require a greater contribution from shocks for filamentary regions than for non-filamentary regions to explain the run of [\\ion{O}{3}]/H$\\alpha$.  In either case, the [\\ion{S}{2}]/[\\ion{N}{2}] ratio is not well fit by the models.  Composite models for UGC 10288 are successful at reproducing the run of [\\ion{S}{2}]/[\\ion{N}{2}] for all but the the highest values of [\\ion{N}{2}]/H$\\alpha$;  however, the run of [\\ion{O}{3}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$ does not show any discernible trend, making it difficult to determine whether or not shocks contribute to the layer's maintenance.  We also examine whether the data can be explained simply by an increase in temperature with $z$ in a pure photo-ionization model without a secondary source of ionization. Runs of [\\ion{S}{2}]/H$\\alpha$, [\\ion{N}{2}]/H$\\alpha$, and [\\ion{S}{2}]/[\\ion{N}{2}] in each of the four galaxies are consistent with such an increase.  However, the runs of [\\ion{O}{3}]/H$\\alpha$ vs. $z$ in NGC 5775 and UGC 10288 require unusually high ionization fractions of O$^{++}$ that can not be explained without invoking a secondary ionization source or at the very least a much higher temperature for the [\\ion{O}{3}]-emitting component than for the [\\ion{S}{2}]- and [\\ion{N}{2}]-emitting component.  An increase in temperature with $z$ is generally more successful at explaining the [\\ion{O}{3}]/H$\\alpha$ run in NGC 891, with the ionization fraction of O$^{++}$ remaining relatively low and nearly constant with $z$. ",
        "introduction": "A vertically extended layer known as the Reynolds Layer or the Warm Ionized Medium (WIM) of the Milky Way contains the majority of the ionized gas in the interstellar medium.  The layer fills about 20\\% of the ISM volume with a local midplane density of $\\sim$0.1 cm$^{-3}$ and a scale height of $\\sim$1 kpc \\markcite{}(Reynolds 1993).  Energetic arguments favor photo-ionization by massive stars in the disk for the layer's maintenance, as roughly 15\\% their ionizing output is sufficient to maintain the layer in the solar neighborhood \\markcite{}(Reynolds 1993).  The Wisconsin H$\\alpha$ Mapper (WHAM; \\markcite{}Reynolds et. al. 1998b) has been instrumental in recent years in characterizing the WIM.  In external galaxies, where the layer is more commonly referred to as Diffuse Ionized Gas (DIG), analogous information has been obtained through the use of narrow-band emission-line imaging, long-slit spectroscopy, and Fabry-Perot observations [e.g. Rand, Kulkarni, \\& Hester 1990 (hereafter RKH); Rand 1996, 1997, 1998; Golla, Dettmar, \\& Domg\\\"{o}rgen 1996; Ferguson, Wyse, \\& Gallagher 1996; Wang, Heckman, \\& Lehnert 1997; Hoopes, Walterbos, \\& Rand 1999; Collins et. al. 2000 (hereafter CRDW)].  Estimates of the diffuse fraction in DIG layers put the ionization requirement closer to 40\\% of the massive stars' ionizing output (e.g. Hoopes et. al. 1996; Ferguson et. al. 1996).  The advantage in observing external edge-on galaxies is the ability to obtain an overview of the entire halo, so that brightness, spatial distribution, and emission line ratios can be better characterized as they vary with $z$.  One of the most important issues regarding these layers is their ionization.  In the WIM, as well as in external DIG layers, [\\ion{S}{2}] $\\lambda6716/$H$\\alpha$ and [\\ion{N}{2}] $\\lambda6583/$H$\\alpha$ are enhanced relative to [\\ion{H}{2}] regions [e.g. Haffner, Reynolds, \\& Tufte 1999 (hereafter HRT); Ferguson, Wyse, \\& Gallagher 1996; Rand 1997].  Such a scenario is predicted by photo-ionization models in which the dilution of the radiation field, measured by the ionization parameter, $U$, increases as photons propagate away from \\ion{H}{2}\\ regions [Domg\\\"{o}rgen \\& Mathis 1994; Sokolowski 1994 (hereafter S94); Bland-Hawthorn et. al. 1997].  Species such as S which are predominantly doubly ionized in \\ion{H}{2}\\ regions are able to recombine to a singly ionized state as the field becomes increasingly dilute.  Such an effect should not be as significant for N or O, which have considerably higher second ionization potentials and thus may not be fully doubly ionized in \\ion{H}{2}\\ regions.  The photo-ionization scenario is not without problems, however.  Firstly, in an increasingly dilute field, photoionization models predict that the ratio [\\ion{S}{2}]/[\\ion{N}{2}] should increase with $z$ due to differing second ionization potentials of S and N.  However, [\\ion{S}{2}]/[\\ion{N}{2}] vs. $z$ has been found to remain nearly constant in the WIM \\markcite{}(HRT) and NGC 891 \\markcite{}(Rand 1998), and in the case of NGC 4631, the variations with $z$ are not as pronounced as those predicted by the models \\markcite{}(Golla et al. 1996).  In addition, the ratio [\\ion{O}{3}] $\\lambda5007/$H$\\alpha$ should decrease with $z$ in a photo-ionization scenario as the field becomes increasingly dilute.  The opposite trend, however, is seen in NGC 891 (Rand 1998), where [\\ion{O}{3}]/H$\\alpha$ increases significantly from the midplane to $z=2$ kpc, as well as in NGC 5775 (Rand 2000; T\\\"{u}llmann et. al. 2000).    [\\ion{O}{3}]/H$\\alpha$ enhancements in DIG relative to \\ion{H}{2}\\ regions are also seen in a number of more face-on galaxies, including M51 \\markcite{}(Wang et. al. 1997).  Finally, photoionization models have trouble explaining values of [\\ion{N}{2}]/H$\\alpha$ and [\\ion{S}{2}]/H$\\alpha$ that are often $\\gtrsim1$ (Wang et. al. 1997; Rand 1998; HRT).  Models of \\markcite{}S94, who took observed line ratios in NGC 891 into consideration, can only explain such ratios by assuming an upper mass cutoff for the stellar IMF of $120 M_{\\sun}$, hardening of the radiation field as it propagates, and a depletion of important gas-phase coolants onto dust grains.  Observations of the weak \\ion{He}{1}\\ $\\lambda5876$ line, however, in the Reynolds Layer (Reynolds \\& Tufte 1995) and in NGC 891 (Rand 1997) each imply a significantly softer spectrum than indicated by the forbidden lines.  One possibility that has been forwarded to explain some of this line ratio behavior is to assume a secondary ionizing source that contributes a small but increasing fraction of the H$\\alpha$ emission with $z$.  Such secondary sources include shocks \\markcite{}(Shull \\& McKee 1979) or turbulent mixing layers (TMLs; \\markcite{}Slavin, Shull, \\& Begelman 1993), each of which may produce large amounts of [\\ion{O}{3}] emission.  Other, less well-studied ionization sources may be relevant in halos, such as ionization from soft X-rays or cosmic rays.  Another approach that has been taken is to assume, instead of the ionization parameter or state changing with $z$, that the line ratio behavior can be explained by the gas temperature increasing with $z$.  This requires a source of additional heating, such as photo-electric heating from dust grains (Reynolds \\& Cox 1992) or dissipation of interstellar turbulence (Minter \\& Balser 1997), that becomes more significant at lower gas densities, but may not require any additional source of ionization.  A temperature increase with $z$ is successful in explaining enhancements in [\\ion{N}{2}]/H$\\alpha$ and [\\ion{S}{2}]/H$\\alpha$ with $z$, as well the constancy in [\\ion{S}{2}]/[\\ion{N}{2}] in the Reynolds Layer up to $z=1.75$ kpc (HRT; Reynolds, Haffner, \\& Tufte 1999).  Such a rise in temperature could also explain an increase in [\\ion{O}{3}]/H$\\alpha$, though data on the run of [\\ion{O}{3}]/H$\\alpha$ vs. $z$ in the WIM are currently lacking.  The goal of this research is to attempt to better understand how DIG halos are energized.  We attempt to determine whether measured runs of line ratios versus $z$ can be explained by: (\\em{a}\\em) photo-ionization with a contribution from shock ionization that increases with height above the midplane, and/or (\\em{b}\\em) only photo-ionization, but with an increase in gas temperature with $z$.  We examine our data in the context of each of these scenarios, realizing that elements of both could be important for understanding the ionization of DIG halos in general.   To this end, we have obtained data for two galaxies: NGC 4302 and UGC 10288.  We also employ data for NGC 5775, already presented by Rand (2000).  For the purpose of assessing the rising temperature scenario, we use data for NGC 891, which has been compared with composite shock/photo-ionization models previously \\markcite{}(Rand 1998).  The observations of these galaxies are discussed in \\S\\ 2.  Line ratio data are presented in \\S\\ 3 along with comparisons to composite shock/photo-ionization and rising temperature models. We briefly summarize our results in \\S\\ 4. ",
        "conclusions": "Deep spectroscopy of four edge-on spiral galaxies has allowed physical conditions within their DIG halos to be probed, and has provided information on possible sources of ionization of the gas.  This study has revealed the following conclusions:  1. The general trend for the runs of line ratios for NGC 4302 and UGC 10288 is very similar to that of NGC 5775 \\markcite{}(Rand 2000) and NGC 891 \\markcite{}(Rand 1998): an increase in [\\ion{N}{2}]/H$\\alpha$ and [\\ion{S}{2}]/H$\\alpha$ with [\\ion{S}{2}]/[\\ion{N}{2}] remaining nearly constant with $z$, though one slit position for UGC 10288 does show evidence for an increase in [\\ion{S}{2}]/[\\ion{N}{2}] with $z$.  Determinations of halo emission scale height have revealed that the more prominent DIG halos also have greater scale height.  This is consistent with models of dynamic halos where greater levels of star formation in the disk can push material farther off the midplane.  2. Pure photo-ionization models do not explain the runs of line ratios in NGC 5775 and UGC 10288.  Composite photo-ionization/shock models can replicate the run of [\\ion{O}{3}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$ in NGC 5775, though the run of [\\ion{S}{2}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$ is still problematic, especially at low-$z$.  In addition, we find that the composite models require a greater percent contribution from shocks in filamentary than non-filamentary regions to match the run of [\\ion{O}{3}]/H$\\alpha$.  A lack of a discernible trend for the [\\ion{O}{3}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$ run in UGC 10288 makes it difficult to determine whether shocks contribute to line emission, though the high values of [\\ion{O}{3}]/H$\\alpha$ in Slit 1 indicate that a shock contribution may be necessary.  A composite model, though, can account for the run of [\\ion{S}{2}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$.  However, if it follows the trend seen in NGC 5775 of more disturbed regions requiring a larger shock contribution, then the lack of a bright, filamentary halo in UGC 10288 may imply that shocks do not play as significant a role in layer's ionization.  The [\\ion{O}{3}]/H$\\alpha$ vs. [\\ion{N}{2}]/H$\\alpha$ run seems to be a key discriminant for shocks.  Fabry-Perot observations of these 3 lines for brighter halos, such as NGC 5775, are now feasible and would provide much more information on the importance of secondary ionizing sources with local morphology and height.  It is important to emphasize that shocks are not the only possible secondary source.  Other sources such as TMLs, X-rays, and ionization by cosmic rays may contribute as well.  3. Another scenario that we consider is a pure photoionization model with an increasing temperature with $z$. By making a few reasonably valid assumptions, the run of gas temperature vs. $z$ has been determined in each of the four galaxies.  We find that an increase in temperature with $z$ can explain the general trend of an increase in the ratios of forbidden lines to Balmer lines with $z$, including an increase in [\\ion{O}{3}]/H$\\alpha$ and the constancy of [\\ion{S}{2}]/[\\ion{N}{2}].  By establishing the run of gas temperature, constraints on ionization fractions of Sulfur, Oxygen, and Hydrogen could then be established.  The [\\ion{O}{3}]/H$\\alpha$ ratios in NGC 5775 and UGC 10288 generally indicate an exceptionally high ionization fraction of O$^{++}$, that invalidates the assumption that N$^{+}$ and H$^{+}$ have equal ionization fractions, indicating that the line ratios cannot be explained by a temperature increase alone.    The runs of line ratios in NGC 891, where the ionization fraction of O$^{++}$ is still somewhat high, and NGC 4302 are better explained by an increase in temperature with $z$, though the absence of [\\ion{O}{3}]$\\lambda$5007 data for NGC 4302 makes a proper assessment difficult.  Finally, [\\ion{O}{1}]$\\lambda$6300, [\\ion{O}{2}]$\\lambda$3727, and [\\ion{O}{3}]$\\lambda$5007 data for a vertically oriented slit would provide additional information on ionization and excitation conditions in these DIG halos."
    },
    "0011/astro-ph0011230_arXiv.txt": {
        "abstract": "GRS 1915+105 has exhibited at least three modes of long-wavelength (infrared to radio) flares.  Class A flares are the bright ($\\sim 1$ Jy) radio events whose apparent superluminal motion in many ways defines the microquasars.  Class B flares have intermediate ($\\sim 50-300$ mJy) flux densities in the IR and radio, and are associated with hard X-ray dipping behaviors.  Class C flares are faint ($\\sim 5-50$ mJy) IR/radio events associated with soft X-ray dips.  We discuss each of these classes and their inter-relationships. ",
        "introduction": "The discovery of superluminal motions in a radio flare from GRS 1915+105 by Mirabel \\& Rodriguez (1994) in many ways launched the field of microquasar astrophysics.  Such flares have been seen since then (e.g. Fender et al., 1999), and always seem to have peak flux densities of $F_{\\nu} \\sim 1$ Jy.  The resolved jets with $v = 0.92 c$ show spectral evolution from optically thick to optically thin, and seem to recur every few days during an ejection episode.  The ejection episodes themselves recur on timescales of a few years.  While Harmon et al. (1997) observed an association between the Class A flares and hard X-ray activity, no simultaneous time-resolved X-ray coverage has yet been obtained. ",
        "conclusions": ""
    },
    "0011/astro-ph0011006_arXiv.txt": {
        "abstract": "We report $R$-- and $V$--band photometry derived from CCD imaging for objects in nine fields in and around the Draco dwarf spheroidal galaxy.  The most distant fields are about 1.3$^\\circ$ from the center.  We use these data to search for Draco stars outside of its measured tidal boundary.  The search involves three methods:  1) Plotting color-magnitude diagrams for individual fields, for sections of fields, and for combined fields and sections.  A color-magnitude diagram can reveal a population of Draco stars by the presence of the expected principal sequences.  2) Measuring field-to-field fluctuations and 3) measuring intra-field fluctuations in the surface density of objects located near the Draco principal sequences in the color-magnitude diagram.  We find evidence for the presence of Draco stars immediately beyond the measured tidal boundary of Draco and place an upper limit on the number of such stars in more distant fields that lie close to the extension of its major axis.  The best evidence is the presence of the Draco principal sequences in the color-magnitude diagram for some combined fields and sections of fields.  The measurements of the field-to-field fluctuations in the stellar surface density confirm this result. ",
        "introduction": "\\label{intro}  It is a long-standing question whether the tidal forces of the Milky Way have played, and continue to play, a significant role in the formation and evolution of its companion dwarf spheroidal (dSph, hereafter) galaxies (e.g., Hodge \\& Michie 1969).  Proposed effects of these forces on a dSph are episodic star formation (Lin \\& Murray 1998), a spuriously large measured mass-to-light ratio (e.g., Kuhn \\& Miller 1989; Bellazzini \\etal\\ 1996; Klessen \\& Kroupa 1998; but also see Pryor 1996; Olszewski 1998; Hirashita \\etal\\ 1999), and structural alterations of the outer regions.  Examples of structural alterations are isophote twisting, truncation of the radial density profile, and the formation of extended streams of tidal debris.  In this article we look for the last of these effects using star counts in two colors to search for stars of the Draco dSph beyond its measured tidal boundary.  Irwin \\& Hatzidimitriou (1995; hereafter IH) list the most accurately measured value of the ``tidal radius'' of Draco: $28.3 \\pm 2.4$~arcmin along the major axis.  They obtained this value by fitting a single-component King (1966) model to their projected density profile derived from star counts made using Palomar Schmidt plates digitized and analyzed using the APM facility (Kibblewhite \\etal\\ 1984). Because there is no fundamental reason why a King model must describe a dSph and, as IH note, the models do not fit the outer parts of their profiles well, this measured tidal radius should not be literally interpreted either as the limiting radius of a dSph or as the radius beyond which stars are unbound due to Galactic tides.  Thus, this tidal radius should only be considered as a parameter in a fitted model. Throughout this article we use the phrase ``tidal boundary'' to mean the elliptical contour centered on Draco with a semi-major axis of 28.3~arcmin, an ellipticity of 0.29, and position angle of 82~degrees (IH).  However, we do not automatically interpret Draco stars beyond this boundary as being unbound from Draco by Galactic tidal forces.  Even though the physical interpretation of the tidal radius resulting from fitting a King model to the projected density profile of a dSph is uncertain, the spatial structure of the outer regions of a dSph is an important source of information about the galaxy.  The density profile of bound stars at large radii is needed to constrain the central density of dark matter in a dSph (Pryor 1994).  Moore (1996) has emphasized that the radius beyond which most of the stars are unbound due to Galactic tidal forces constrains the total mass of a dSph and, hence, its total amount of dark matter.  Johnston \\etal\\ (1999) identify this radius with that where the slope of the projected density profile in a log-log plot abruptly becomes shallower, based on numerical simulations of dSphs interacting with the tidal field of the Galaxy.  The unbound stars also should show velocity gradients and a velocity dispersion increasing with increasing distance from the dSph (Oh \\etal\\ 1995; Piatek \\& Pryor 1995; Johnston \\etal\\ 1999). Observing these kinematic signatures is necessary to confirm that these stars are unbound from the dSph by Galactic tidal forces, since other mechanisms could produce an extended bound population with a shallow density profile at large radii.  An example is a halo of stars that formed from gas pushed to large radii by supernovae and stellar winds.  In this article, we report $R$-- and $V$--band photometry for objects in nine fields in and around Draco and the results of using these data to search for Draco stars beyond its measured tidal boundary.  Section 2 describes the data acquisition and reduction.  Section 3 presents color-magnitude diagrams for our fields and discusses the classification of an object based on its location in the diagram and the morphology of its image.  Section~4 describes our methods of searching for Draco stars and the results they yield. Section 5 summarizes our results and compares them to the results from other similar studies. ",
        "conclusions": "\\label{discussion}  The presence of principal sequences in the CMD (Figures 12~(a) and 13(a)) for the regions of the E1, W1, and SE1 fields beyond the tidal boundary measured by IH demonstrate the presence of Draco stars in at least some of these regions.  The value of $\\Sigma_{\\rm in}$ for the region of W1 beyond the measured tidal boundary confirms that Draco stars are in this region.  We think that the CMDs shown in Figures~12(c) and 13(c) provide evidence (primarily in the visibility of the HB) that the Draco stars extend into the more distant fields E2 and W2.  However, we acknowledge that others may disagree because the principal sequences are present only at the level of the noise.  The field-to-field scatter in the values of $\\Sigma_{\\rm in}$ for the more distant fields is comparable to that of $\\Sigma_{\\rm out}$ and so the values of $\\Sigma_{\\rm in}$ can neither prove nor disprove the presence of Draco stars at these larger radii.  A confirmation of the membership of these stars must come from other techniques, such as measurements of radial velocity, proper motion, metallicity, or luminosity.  If there are Draco stars in our E2 and W2 fields, this implies that the surface density of Draco decreases with increasing radius much more slowly beyond a radius of about 30~arcmin than between 10 and 20~arcmin.  Models of the tidal stripping of stars from a dSph produce such a break in the profile at the approximate radius where the stars change from being mostly bound to mostly unbound (e.g., Johnston et al. 1999).  If this interpretation is correct, the stars at large radii should have a mean velocity above that of Draco on one side and below on the other (Oh \\etal\\ 1995; Piatek \\& Pryor 1995; Klessen \\& Kroupa 1998; Johnston et al. 1999).  Until such measurements exist, alternative explanations are also viable.  For example, the creation of a spatially extended population of bound stars either during the formation of the dSph or later when gas was driven to large radii by stellar winds or supernovae.  Gas is seen in large shells around some dwarf irregular galaxies (Puche \\& Westpfahl 1994).  The greater visibility of the principal sequences in the CMDs of Figures~12(c) and 13(c) compared to those of Figures~12(d) and 13(d) indicates that any Draco stars present at large radii concentrate mainly to the east and west of the galaxy.  This suggests that the stars beyond the measured tidal boundary are along extensions of the major axis, but demonstrating this requires data with more complete azimuthal coverage.  We find that searching for Draco stars through fluctuations in the surface densities of objects in fields where the expected number of these stars is small compared to the numbers of foreground field stars and background galaxies is subject to significant uncertainties.  The left-hand panels of Figure~15 show that galaxies produce fluctuations larger than expected from Poisson noise over a broad range of limiting magnitudes, thus potentially masking fluctuations due to Draco stars. Removing objects based on the morphology of their images can greatly reduce these fluctuations (see the right-hand panels of Figure~15). However, this procedure cannot be completely effective, particularly when applied to objects with low signal-to-noise ratios, and can introduce spurious fluctuations if the focus or seeing varies among fields.  We estimate the surface density of Draco stars beyond the measured tidal boundary using the $\\Sigma_{in}$ values from the third block of Table~7 and the equivalent value for the W1 field given in the text. The most reliable estimate is for this region of the W1 field.  The difference between the $\\Sigma_{in}$ for this region and the mean value for the SE1, N2, E2, S2, SW2, and S2 fields, given in Table~7, yields $0.30 \\pm 0.10$~stars~arcmin$^{-2}$, while using the lowest $\\Sigma_{in}$, that for the SE1 field, instead of the mean yields $0.39 \\pm 0.11$~stars~arcmin$^{-2}$.  The comparison of these values with those of IH and with theoretical models requires that they be expressed as a fraction of the central surface density of Draco.  To calculate a central surface density of Draco, we derived its density profile along the major axis in the C0 field by counting in elliptical annuli those objects that contribute to the $\\Sigma_{in}$ listed in the third block of Table~7.  The ellipticity and position angle of these annuli are from IH.  The greater incompleteness of the C0 field (see Table~5) was approximately corrected by multiplying the number of objects between $R = 22.1$ and 22.6 (38\\% of the total number) by 0.95/0.80 = 1.19.  The resulting profile agrees well with that of IH and yields a central surface density of $20.0 \\pm 1.4$~stars~arcmin$^{-2}$.  The Poisson uncertainties in the binned densities account for most of the uncertainty in this value.  We will present the radial density profile for Draco based on the data from all of our fields elsewhere (Piatek \\etal\\ 2000).  Our estimate of the surface density of Draco stars in the W1 field beyond the measured tidal boundary is thus $0.015 \\pm 0.005$ to $0.020 \\pm 0.006$ of the central density.  This estimate is within the range of values at this radius shown in Figure~2 of IH.  An estimate of the surface density of Draco stars in a field at larger radii is the difference between its value of $\\Sigma_{in}$ in the third block of Table~7 and the lowest value of $\\Sigma_{in}$, that for the S2 field --- which is assumed to be the background.  The largest difference, for the SW2 field, is $0.18 \\pm 0.08$~stars~arcmin$^{-2}$ or $0.0090 \\pm 0.0039$ of the central surface density.  We emphasize that this estimate is suspect both because the SW2 field contains a rich cluster of galaxies which contaminates $\\Sigma_{in}$ and because the differences are not statistically significant (see the discussion in Sec.~\\ref{fluctsearch}).  The second largest difference is for the E2 field: $0.17 \\pm 0.08$~stars~arcmin$^{-2}$ or $0.0083 \\pm 0.0039$ of the central surface density.  These values are consistent with the lowest point of the profile shown in Figure~2 of IH (though this is only the lowest point that is above their assumed background).  If there are Draco stars in the E2 and W2 fields and not in the N2, S2, and SW2 fields, as suggested by Figures~12(c) and (d) and 13(c) and (d), then the above value for the E2 field is probably an upper limit on their surface density.  However, we again emphasize that the evidence for field-to-field fluctuations in the surface density of Draco stars at large radii is weak (see Sec.~\\ref{fluctsearch}).  Several other groups have reported discoveries of stars near or beyond the tidal boundary of the Sextans, Carina, and Sagittarius dSphs.  Gould \\etal\\ (1992) identified at least five faint metal-poor dwarf stars, apparently all at a common distance of about 100~kpc, in their CT1 field.  This field is about 100~arcmin away from the center of Sextans along the major axis.  The properties of these stars are consistent with membership in Sextans and Gould \\etal\\ (1992) estimate a surface density, based on these stars, of about 0.01 of the central value for Sextans.  The IH value for the major-axis tidal radius of Sextans is $160 \\pm 50$~arcmin, implying that the five stars are within the tidal boundary.  The surface density found by Gould \\etal\\ (1992) is consistent with the IH profile for Sextans.  Kuhn et al.\\ (1996; hereafter KSH) found evidence for Carina stars in fields up to $2^{\\circ}$ away from the center of this galaxy along the extension of its major axis.  Their approach is similar to our method of searching for field-to-field fluctuations in $\\Sigma_{\\rm in}$, except that it subtracts an estimate of the contribution from non-member objects.  They use least squares to estimate the contribution to the CMD of each field from both the CMD of a background field, containing non-member stars and galaxies, and a CMD which is the difference between the CMD of a field centered on Carina and the background CMD, intended to contain only Carina stars.  Out of the eight fields along the extension of the major axis and not centered on Carina, seven have surface densities of Carina stars larger than those of the two fields 1$^\\circ$ on either side of the center along the extension of the minor axis, which were taken to represent the background.  The two nearest of the eight fields have the highest measured contribution from the Carina CMD.  The center of the nearest field is slightly within the tidal boundary measured by IH, so the presence of Carina stars in this field is expected and, as KSH note, is consistent with the profile measured by IH.  The second nearest field is just outside of the measured tidal boundary of Carina, so the inferred presence of Carina stars there is reminiscent of our evidence for Draco stars in the regions of the E1 and W1 fields outside of the tidal boundary of Draco.  Our values for $\\Sigma_{\\rm in}$ for fields about 1$^\\circ$ from the center of Draco, shown in the left-most panel of the bottom row of Figure~16, display less scatter than do the values for the contribution of the Carina CMD to the CMDs of the KSH fields. This is the case even after excluding the two innermost of their fields.  The scatter is closer to what we see before removing galaxies (see the left-most panel of the top row of Figure~16).  KSH do not remove galaxies from their sample.  The uncertainty for the Carina contribution estimated by KSH comes from the scatter of the values from the four quadrants of the field, as do our error bars in the bottom row of Figure~16.  Their uncertainties vary more between fields than do ours.  This may be due to either or both of the following:  the clustering of galaxies on angular scales comparable to the size of the field and the small numbers of objects in each bin of their least-squares fit.  Because of the above concerns about galaxies and the variable uncertainties, we think that evidence for Carina stars 1$^\\circ$ or more from the center is weaker than KSH claim and is probably no stronger than our evidence for Draco stars in the E2 and W2 fields.  Majewski~\\etal\\ (2000; hereafter M2K) identified a ``substantial extratidal population from Carina'' using Washington-band and DDO51-band photometry.  With this technique it is possible to distinguish metal-poor giants from the more numerous disk and halo dwarf stars (Geisler 1984) to a degree that depends on photometric accuracy.  Based on the size of this population beyond the tidal boundary measured by IH, M2K estimated that Carina loses about 27\\% of its mass per Gyr.  Spectroscopy of three bright giant candidates outside of the tidal boundary confirmed that they have velocities consistent with that of Carina.  Notwithstanding this confirmation, Morrison~\\etal\\ (2000) argued that photometric errors of 0.1~mag, the limit adopted by M2K, cause dwarfs from the Galactic disk and halo to contaminate significantly the population of fainter candidate giants. The photometric and spectroscopic data of M2K demonstrate that there are Carina stars beyond its measured tidal boundary.  However, the number of these stars and their distribution in azimuth and radius remain uncertain because of the incomplete areal coverage and variable limiting magnitudes of the M2K fields, and the concerns raised by Morrison~\\etal\\ (2000).  Since the discovery of the Sagittarius dSph (Ibata \\etal\\ 1994), stars associated with this galaxy have been found at ever increasing angular distances from its center.  Mateo \\etal\\ (1996) and Fahlman \\etal\\ (1996) found main-sequence stars of Sagittarius about 10$^\\circ$ from the center.  Mateo \\etal\\ (1996) also found RR~Lyrae stars associated with Sagittarius in their field, as did Alard (1996) and Alcock \\etal\\ (1997) in fields 8$^\\circ$ to 11$^\\circ$ from the center on the opposite side.  Ibata \\etal\\ (1997)  detected main-sequence stars in a $22^{\\circ} \\times 8^{\\circ}$ region on the side of Sagittarius away from the Galactic plane.  They infer a minor axis limiting radius of 4$^\\circ$ and a major to minor axis ratio of 3:1, implying a major axis limiting radius of 12$^\\circ$.  However, Mateo \\etal\\ (1998) detected main-sequence stars over the range of 10$^\\circ$ to 34$^\\circ$ from the center of Sagittarius.  Their measured surface density profile along the major axis has a break at a radius of about 20$^\\circ$.  Majewski \\etal\\ (1999) report a possible detection of Sagittarius stars at similar angular distances.  Recently, a survey for carbon stars in the Galactic halo (Totten \\& Irwin 1998; Totten \\etal\\ 2000) has found giant carbon stars distributed on a great circle passing through the position of Sagittarius.  Ibata \\etal\\ (2000a) interpret these stars as a stream of tidal debris extending completely around the orbit of Sagittarius.  Ibata \\etal\\ (2000b) also argue that halo substructure found in the Sloan Digital Sky Survey (Yanny \\etal\\ 2000; Ivezi\\'{c} \\etal\\ 2000) is tidal debris from Sagittarius.  The existence of stars beyond the tidal boundary of Sagittarius is established and they are very likely due to tidal stripping.  Is this the case for the Galactic dSphs at larger galactocentric distance, for which tides are currently less important?  Our study and others, particularly that of Majewski \\etal\\ (2000) for Carina, show that several dSphs extend beyond the tidal boundaries measured by IH. However, these stars have not been found to extend to as large a distance along the major axis as for Sagittarius, so alternative explanations remain viable.  The most mundane explanation is that these stars are within the tidal boundary of their dSph.  This boundary is uncertain because there is no fundamental reason why a King model (King 1966) should correctly describe a dSph at all radii. For example, a dSph could be surrounded by an extended population of bound stars formed when gas in the galaxy was expelled by supernovae and stellar winds.  Convincing evidence for tidal stripping would be the existence of a very extended distribution of stars along some axis, particularly if supported by the presence of a velocity gradient in these stars (Oh \\etal\\ 1995; Piatek \\& Pryor 1995; Johnston \\etal\\ 1995).  Though this axis must be along the orbit, it is not clear that this must be the major axis of the bound stars since the galaxy might have formed with the major axis not in the orbital plane.  Indeed, an alignment between these two could argue for a strong tidal interaction, though this might also be a result of the dSph forming in the vicinity of the Galaxy. Extensions of the work presented in this paper that could provide better evidence for tidal stripping are deeper two-color photometry with good seeing -- preferably followed up with radial velocity or proper motion measurements, searches for RR~Lyrae variables, and three-color photometry that is able to distinguish metal-poor giants from disk and halo dwarfs."
    },
    "0011/astro-ph0011376_arXiv.txt": {
        "abstract": "High-resolution N-body simulations of four popular Cold Dark Matter cosmologies (LCDM, OCDM, QCDM, and tilted SCDM), each containing $\\sim 10^{5}$ clusters of galaxies of mass $M_{1.5}>5\\times 10^{13}h^{-1}M_{\\odot }$ in a Gpc$^{3}$ volume, are used to determine the evolution of the cluster mass function from $z=3$ to $z=0$. The large volume and high resolution of these simulations allow an accurate measure of the evolution of cosmologically important (but rare) massive clusters at high redshift. The simulated mass function is presented for cluster masses within several radii typically used  observationally ($R=$ 0.5, 1.0, and 1.5 $h^-1$Mpc, both comoving and physical) in order to enable direct comparison with current and future observations. The simulated evolution is compared with current observations of massive clusters at redshifts $0.3 \\lesssim z \\lesssim 0.8$. The $\\Omega_m$=1 tilted SCDM model, which exhibits very rapid evolution of the cluster abundance, produces too few clusters at $z\\gtrsim 0.3$ and no massive clusters at $z\\gtrsim 0.5$, in stark contradiction with observations. The $\\Omega_m$=0.3 models--- LCDM, OCDM, and QCDM--- all exhibit considerably weaker evolution and are consistent with current data. Among these low density models, OCDM evolves the least.  These trends are enhanced at high redshift and can be used to discriminate between flat and open low density models.  The simulated mass functions are compared with  the Press-Schechter approximation. Standard Press-Schechter predicts too many low mass clusters at $z=0$, and too few clusters at higher redshift. We modify the approximation by a simple parameterization of the density contrast threshold for collapse, which has a redshift dependence. This modified Press-Schechter approximation provides a good fit to the simulated mass functions. ",
        "introduction": "\\label{secintro}  The local abundance of clusters of galaxies places a powerful constraint on cosmological parameters: $\\sigma _{8}\\Omega _{m}^{0.5}$ $\\simeq $ 0.5, where $\\sigma _{8}$ is the rms mass fluctuation on an 8$h^{-1}$Mpc scale and $\\Omega_{m}$ is the present cosmological density parameter (\\cite{HA91}, \\cite{BC92}, \\cite{WEF93}, \\cite{ECF96}, \\cite{VL96}, \\cite{KS97}, \\cite{Pen98}). This constraint, while powerful, is degenerate in $\\Omega _{m}$ and $\\sigma _{8}$. The evolution of cluster abundance with redshift, especially for massive clusters, breaks this degeneracy (e.g., \\cite{PDJ89}, Oukbir \\& Blanchard (1992, 1997), \\cite{ECF96},  \\cite{VL96}, \\cite{BFC97}, \\cite{FBC97}, \\cite{CYE97}, Henry (1997, 2000), \\cite{BF98}, \\cite{ECFH98}, \\cite{DV99}). The evolution of high-mass clusters is strong in $\\Omega _{m}$ = 1, low-$\\sigma _{8}$ Gaussian models, but is much weaker in low-density, $% \\sigma _{8}$ $\\sim$1 models. Therefore, the evolution of cluster abundance provides one of the most powerful methods to determine both $% \\Omega _{m}$ and $\\sigma _{8}$. Various analyses using this method have shown that the relatively high abundance of observed clusters to $z\\lesssim 0.8$ indicates that the mass-density of the universe is low: $\\Omega_{m} \\sim 0.2$ - 0.5, and $\\sigma _{8} \\sim$ 1 (see references above). However, the total number of clusters currently studied at high redshifts is still small, and the model comparisons are generally based on the Press-Schechter (1974) approximation which,  while surprisingly accurate, is known to have biases. Direct comparisons with simulations have also been used, but these are generally too small to find the rare high-mass clusters, especially at large redshifts. New observations are currently underway to determine more precisely the evolution of cluster abundance using optical, X-ray, Sunyaev-Zel'dovich, and gravitational lensing surveys. At the same time, larger and more accurate cosmological simulations are needed in order to accurately determine the expected mass function of clusters of galaxies and its evolution with time, in order to allow for proper comparison with the upcoming observations.  In this paper we present gigaparsec-scale simulations of four popular CDM models which allow an accurate determination of the evolution of the cluster mass function from redshift $z=3$ to $z=0$, including the rare and cosmologically powerful massive clusters. The rapid increase of available computing power and new algorithms which run efficiently on parallel machines make possible a combination of large simulation volume and high resolution. For comparison, the models presented here are of a larger volume than the \\cite{GBQTBKL99} simulations, and attain considerably higher resolution than the Virgo ``Hubble Volume'' simulations \\citep{Evrard98,JFWCCEY00}. Furthermore, we are able to cover several different models, including the first large volume, high resolution simulation of tilted $\\Omega_{m}=1$ as well as a quintessence model.   The aim of this paper is to provide the foundation for comparisons with the upcoming observations. Thus the evolution of the mass function is presented in a manner that can be directly compared with observations, using cluster masses within a variety of specific radii typically used in observations, rather than using the not so easily observable virial mass. We discuss the simulations and the method of cluster selection in $\\S$\\ref{seccreate}. The evolution of the resulting cluster mass function is presented in $\\S$\\ref{secevol}. Comparison with the Press-Schechter approximation is made in $\\S$\\ref{secpscomp}, and effects of resolution are discussed in $\\S$\\ref{secreseff}. Comparisons with current observations are  shown in $\\S$\\ref{secobscomp}, and conclusions are summarized in $\\S$\\ref{secconcl}. ",
        "conclusions": "\\label{secconcl}  Gigaparsec scale high-resolution N-body simulations are used to determine the cluster mass function from $z=3$ to $z=0$ for four CDM models: LCDM, OCDM, QCDM, and tilted SCDM. These large, high-resolution simulations, with 10$^{5}$ clusters of mass M$_{1.5} \\geq 5\\times 10^{13}h^{-1}M_{\\odot }$ at the present day, allow an accurate determination of the MF evolution over this redshift range.  The evolution of the  mass function is presented for cluster masses within several radii typically used in observations: 0.5, 1, and 1.5$h^{-1}$% Mpc, both comoving and physical; this will enable direct comparisons with different observations. The evolution of the MF for clusters of a given mass within a physical radius is found to be considerably weaker than for mass within a comoving radius; this is due to the relatively larger physical radii at high redshift, which means that more numerous poorer clusters are included. After virialization at $z \\approx 2$ in the $\\Omega_{m}=0.3$ models, there is relatively little evolution in physical coordinates. Because of the large differences in the resulting evolution, it is essential that the cluster evolution be compared using the same simulated cluster mass--- within the same radius--- as observed.  We compare the simulation results with the Press-Schechter approximation; the latter predicts too many low mass clusters at $z=0$ and too few clusters at higher redshifts. We provide refined parameters for this approximation that best fit the simulations as a function of redshift and mass. With the chosen fits to the overdensity parameter $\\delta_c$, the agreement with the simulation MF is better than 10\\% for $z\\leq 2$ in the three $\\Omega_m=0.3$ models and $z\\leq 1$ in SCDM.  We compare the simulation results with current observations of the most massive distant clusters observed to  $z\\simeq$0.8. We show that Gaussian $\\Omega _{m}$=1 models such as tilted SCDM, which predict extremely rapid evolution from  $z\\sim$1 to  $z\\sim$0, are ruled out at a high significance level; they produce no massive clusters at high redshifts (10$^{-5}$ clusters per the observed volume), in contradiction with observations. Showing much weaker evolution, the $\\Omega _{m}$ = 0.3 models--- LCDM, OCDM, and QCDM--- are all consistent with the data. These results provide one of the most powerful constraints on the mass density of the universe. Upcoming surveys of distant clusters should be able to provide further discrimination between flat (LCDM or QCDM) and open (OCDM) low density cosmologies, since the latter shows relatively weaker evolution.    \\noindent This research was supported by NSF Grants AST-9318185 and AST-9803137 (under Subgrant 99-184), and the NCSA Grand Challenge Computational Cosmology Partnership under NSF Cooperative Agreement ACI-9619019, PACI Subaward 766."
    },
    "0011/astro-ph0011140_arXiv.txt": {
        "abstract": "{ We present the first results of an {\\em XMM-Newton EPIC} survey of NGC 2516, a southern low-metallicity open cluster with an age close to the Pleiades.  The attained limiting sensitivity is of $\\sim 2.4\\times 10^{-15}$ erg~sec$^{-1}$~cm$^{-2}$ in the 0.1--4.0 keV bandpass.  This has been achieved by summing the data of the MOS and PN cameras of two distinct observations for a total exposure time of $\\sim$ 33 ks and by analyzing the summed data set with the wavelet detection algorithm developed at Osservatorio Astronomico di Palermo (OAPA) that has yielded over 200 X-ray detections. Using data just from a single exposure or from a single camera would have reduced by a factor 2--4 our limiting sensitivity and would have resulted in 25--40\\% less X-ray detections.  To date, 129 detections have as counterparts one or more of the 540 photometrically selected cluster members in the surveyed region, for a total of 147 likely detected members, with unique identification in 112 cases.  We derive the X-ray luminosity functions (XLF) of NGC~2516 members of different spectral types and compare them with those of the more metal rich, approximately coeval Pleiades cluster, finding the NGC~2516 photometrically selected dG and dK stars less luminous than the Pleiades. The XLFs of the NGC~2516 and of the Pleiades dM stars are indistinguishable.  We compare the {\\em XMM-Newton} results with those recently obtained with {\\em Chandra}. ",
        "introduction": "NGC~2516 has been chosen as an {\\em XMM-Newton} calibration target in order to ``bore sight'' the alignment of the X-ray telescope as well as the relative positions of the EPIC cameras CCDs since it is a region rich of X-ray emitting young stars. Apart from this role in the early phase of the {\\em XMM-Newton} mission, NGC~2516 plays an important role in the study of stellar physics. It is a southern, relatively nearby ($\\sim$ 387 pc; \\cite{JTP97}), open cluster with metallicity below solar ([Fe/H]~=~ $-$0.32$\\pm$0.06, \\cite{JTP97}; $-$0.42, \\cite{Cam85}; $-$0.18$\\pm$0.08, \\cite{JJT98}), affected by moderate extinction ($E(B-V)~=~0.12$), and an age ($\\sim$ 10$^8$ yr) similar to that of the Pleiades.  Because of its metallicity, it plays a crucial role in exploring the relation between stellar structure, metallicity, and coronal emission level. Such a relation is expected to exist since metallicity affects the depth of the convection zone that in turn should affect the efficiency of dynamo action, of which X-ray luminosity is a proxy (cf. discussion in Micela et al., 2000). Indeed Jeffries et al. (1997), based on {\\em ROSAT-PSPC} data, have found that the brightest X-ray luminous dG NGC~2516 members are less luminous than the brightest X-ray luminous solar-metallicity, similar age, dG in the Pleiades. More recently Harnden et al. (2000), based on {\\em Chandra} data, have found that the median X-ray luminosity of NGC~2516 dF stars is higher than that of the Pleiades, while there is an indication of the NGC~2516 dG and dK being less X-ray luminous than the Pleiades (albeit this result can be subject to uncertainties in assessing member status).  The availability of deep {\\em XMM-Newton} data allow us to further investigate this matter.  Our paper is organized as follows: in \\S 2 we will describe the X-ray data and analysis, in \\S 3 we will describe the results we have obtained and in \\S 4 we will summarize and discuss our findings. ",
        "conclusions": "NGC 2516 was chosen as a calibration target in order to ``bore sight'' the alignment of the X-ray telescope with the {\\em XMM-Newton} detectors, however this cluster is of particular scientific interest since, being metal-poor with respect to the Sun, it allows us to explore the effect of metallicity on coronal emission level.  Using a 33 ks long EPIC data set obtained by summing the data, all taken with the thick filter, of the three distinct EPIC cameras of two distinct co-pointed observations we have reached a limiting sensitivity of $\\sim 2.35\\times 10^{-15}$ erg~sec$^{-1}$~cm$^{-2}$ and have detected 208 sources. Using only data from a single observation and/or from the MOS or PN cameras alone we would have reached a factor 2--4 worst sensitivity and would have detected substantially less sources. The attained limiting sensitivity has allowed us to derive luminosity functions of NGC~2516 members with a larger number of detected members with respect to the {\\em Chandra} survey (\\cite{H++00}). Together with better photometry made available to us before publication (Jeffries, Thurston and Hambly, in preparation), the {\\em XMM-Newton} calibration observations have enabled us to investigate the late-type population of cluster members and to explore their coronal emission properties.  dG and dK stars in NGC~2516 tend to be less luminous than the same spectral type Pleiades members reinforcing the indication coming from the recent analysis of {\\em Chandra} observations (\\cite{H++00}). NGC~2516 dM members have the same coronal emission level of their Pleiades counterparts, confirming the ROSAT HRI-based suggestion of \\cite{MSJ00} that the activity level for dM stars is insensitive to a change of a factor of two in stellar metallicity. On the other hand, we find that the median $\\log(L_{\\rm X})$ value for NGC~2516 dF-type stars when evaluated on the entire EPIC FOV is {\\em lower} than that of the Pleiades, while it has been found {\\em higher} than that of the Pleiades analyzing {\\em Chandra} data. Indeed when we restrict the analysis to the dF stars in the {\\em Chandra} survey FOV we recover the {\\em Chandra} result. Since the {\\em Chandra} FOV is offset pointed with respect to the {\\em XMM-Newton} one, we do not think that the above findings can be explained by considering that we have preferentially included members surveyed in the lower sensitivity outer part of the {\\em XMM-Newton} FOV where we would expect a lower detection percentage. A more viable explanation is that the true members are spatially concentrated, whereas the contaminating non-members are uniformly distributed. In such a picture the {\\em Chandra} observation is pointed where the cluster members are concentrated while the present {\\em XMM-Newton} survey is pointed in a less central cluster region, explaining our finding. However other explanations are possible and further investigations are required.  Lacking astrometric proper motion studies and extensive radial velocity surveys of NGC~2516 stars we have been forced to base our study on photometrically selected members. Such selection is subject to contamination from field stars that by chance have the same magnitude and colors as true NGC~2516 members. Based on the very good quality of the photometry we have used, clearly shown by the CMDs in Fig.~\\ref{fig:CMD}, we think that contamination should not be too severe (but it could be as large as 15--20\\%). However, it is clear that to be able to draw firmer conclusions regarding metallicity effects on coronal emission it is crucial that an improved list of members be obtained complementing available photometric data with high precision proper motion studies (unfortunately NGC~2516 average motion is only marginally different from the solar reflected one) and with an extensive campaign of radial velocity measurements."
    },
    "0011/hep-ph0011019_arXiv.txt": {
        "abstract": "The latent heat released during the expansion of bubbles in the electroweak phase transition reheats the plasma and causes the bubble growth to slow down. This decrease of the bubble wall velocity affects the result of electroweak baryogenesis. Since the efficiency of baryogenesis peaks for a wall velocity  $\\sim 10^{-2}$, the resulting baryon asymmetry can either be enhanced or suppressed, depending on the initial value of the wall velocity. We calculate the evolution of the phase transition taking into account the release of latent heat. We find that, although in the SM the baryon production is enhanced by this effect, in the MSSM it causes a suppression to the final baryon asymmetry. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011255_arXiv.txt": {
        "abstract": "We present some results from our {\\it HST} archival image study of 71 QSO host galaxies. The objects are selected to have $z\\leq0.46$ and total absolute magnitude $M_V\\leq -23$ in our adopted cosmology ($H_0=50$ km s$^{-1}$ Mpc$^{-1}$, $q_0=0.5$, $\\Lambda=0$). The aim of this initial study is to investigate the composition of the sample with respect to host morphology and radio loudness, as well as derive the QSO host galaxy luminosity function. We have analyzed available WFPC2 images in $R$ or $I$ band ($U$ in one case), using a uniform set of procedures.  The host galaxies span a narrow range of luminosities and are exceptionally bright, much more so than normal galaxies, usually $L>L^*_V$.  The QSOs are almost equally divided among three subclasses: radio-loud QSOs with elliptical hosts, radio-quiet QSOs with elliptical hosts, and radio-quiet QSOs with spiral hosts. Radio-loud QSOs with spiral hosts are extremely rare.  Using a weighting procedure, we derive the combined luminosity function of QSO host galaxies.  We find that the luminosity function of QSO hosts differs in shape from that of normal galaxies but that they coincide at the highest luminosities.  The ratio of the number of quasar hosts to the number of normal galaxies at a luminosity $L_V$ is $\\mathcal{R}=(L_V/{11.48L^*_V})^{2.46}$, where $L^*_V$ corresponds to $M^*_V=-22.35$, and a QSO is defined to be an object with total nuclear plus host light $M_V\\leq-23$. This ratio can be interpreted as the probability that a galaxy with luminosity $L_V$ will host a QSO at redshift $z\\approx0.26$. ",
        "introduction": "\\label{sec:intro}  Much has been learned about the properties of QSO host galaxies since they were first imaged almost three decades ago (Kristian~1973). Early results include establishing a positive correlation between host and nuclear QSO luminosities (Hutchings, Crampton, \\& Campbell~1984) and indications of a morphological difference between radio-loud and radio-quiet QSOs, with the former more likely to be in elliptical hosts and the latter in spiral hosts (Malkan, Margon, \\& Chanan~1984). Boroson, Persson, \\& Oke~(1985), as well as Stockton \\& MacKenty~(1987) examine hosts to classify them spectroscopically and in the context of their nuclear emissions.  Working in the near-infrared, where the luminosity contrast is more favorable to the host galaxy, Dunlop et al.~(1993) show that QSO hosts are typically drawn from the bright end of the galaxy luminosity function (in agreement with Hutchings et al.~1984).  McLeod \\& Rieke~(1994a,b), also using near-infrared data, find that hosts of radio-quiet QSOs are typically represented by an exponential (spiral disk) light profile (in agreement with Malkan et al.~1984), and that high-luminosity QSOs generally have brighter hosts than low-luminosity QSOs (in agreement with Hutchings et al.~1984).  High-resolution space-based images taken with the {\\it Hubble Space Telescope} ({\\it HST}) make observing the host galaxy much easier. The first, and to date one of the largest, systematic {\\it HST} studies of QSO hosts is by Bahcall, Kirhakos, \\& Saxe~(1997), who study 20 of the most luminous nearby QSOs.  They can discern the morphology of the hosts, and they discover that, while radio-loud QSOs are found only in ellipticals or interacting systems, radio-quiet QSOs can be in ellipticals, spirals or interacting systems.  They also find that QSO hosts do not follow a Schechter~(1976) luminosity function and are instead found at systematically high luminosities.  More recent studies strengthen and expand on these results.  McLure et al.~(1999) confirm that QSO hosts are generally luminous, and also determine that, even for radio-quiet QSOs, the hosts are often ellipticals or bulge-dominated.  Furthermore, elliptical hosts appear to follow the same luminosity-surface brightness relation as field elliptical galaxies (Hamabe \\& Kormendy~1987).  Other recent studies include Nolan et al.~(2001), who discuss QSO host ages, and Kukula et al.~(2001), who study host evolution from redshifts of $z \\approx 2$ to the local universe.  In this paper we focus on the luminosity distribution of a large sample of QSO hosts observed with the Wide Field and Planetary Camera 2 (WFPC2) aboard {\\it HST} (\\S\\ref{sec:sample}).  We have collected and reanalyzed wide-band archival images of 71 QSOs with $ M_V \\leq -23$ mag (total nuclear $+$ host light) and redshifts $ 0.06 \\leq z \\leq 0.46 $.  We have taken an inclusive approach in our sample selection, imposing no additional selection criteria on the QSOs besides those of total absolute magnitude and redshift, while some of the previous work on QSO hosts has focused on specific classes of QSOs: radio loud (Lehnert et al.~1999), intrinsically very bright (Bahcall et al.~1997), and so on.  For each we have subtracted the nuclear light component using two-dimensional image fits and have derived the luminosity and size of the underlying host galaxy by fitting both an $ r^{1/4} $ and an exponential light profile (\\S\\ref{sec:measurements}).  Given the total number of objects considered, more than triple that of previous studies, we can effectively sample the general QSO population for redshifts $ z \\leq 0.46 $, and derive a global luminosity function for their host galaxies (\\S\\ref{sec:lumfun}) that is not grossly affected by selection criteria.  This luminosity function is compared with that of normal galaxies, and selection effects/biases and other issues are discussed (\\S\\ref{sec:discussion}). Conclusions are then summarized (\\S\\ref{sec:conclusions}).  Throughout this paper, we adopt a Friedman cosmology with $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0.5$, and $\\Lambda=0$.  We have converted the results of other researchers to this cosmology when comparing our results to theirs.  We confirm previous results that host galaxies of QSOs are significantly more luminous than typical luminous $L^*_V$ galaxies, where $L^*_V$ is the ``knee'' in the Schechter (1976) luminosity function.  We also consider the relationship between host morphology and QSO radio loudness.  Spiral hosts are, on average, nearly as luminous as elliptical hosts, while hosts of radio-loud QSOs are, on average, about 0.5 magnitudes brighter than hosts of radio-quiet QSOs. Subject to systematic uncertainties in normalization procedures, we find that the combined low-redshift QSO host luminosity function has a very different shape from that of normal galaxies but that they coincide at the highest luminosities.  In approximate terms, at redshift $z\\approx0.26$ the ratio of the number of these QSO hosts to the number of normal galaxies of luminosity $L_V$ is $\\mathcal{R} \\approx (L_V/{11.48L^*_V})^{2.46}$, where $L^*_V$ corresponds to $M_V^*=-22.35$. ",
        "conclusions": "Here we elaborate on some of the results of this work. A more thorough discussion will be made elsewhere when we consider the other properties of the sample in detail.  \\subsection {Comparison of the QSO Host and Normal Galaxy Luminosity Functions}\\label{sec:normgal}   To compare our QSO host galaxy luminosity function (\\S\\ref{sec:lf}) with that of normal galaxies, we use the normal galaxy luminosity function of Metcalfe et al.~(1998). The Schechter~(1976) luminosity function parameters that describe their V-band luminosity function in our cosmology are $\\alpha=-1.2$, $M^*_V = -22.35$, and $\\Phi^* = 8.5\\times10^{-4}$ Mpc$^{-3}$.  The host luminosity function lies below that of normal galaxies, as shown.  The most relevant uncertainty in the normal galaxy luminosity function is at the bright end, where it is less well constrained due to the dearth of luminous galaxies in surveys.  The least luminous QSO host used in our analysis is relatively luminous, with $M_V = -21.7$, and the median of our QSO host luminosity function is at $M_V \\approx -22.95$, twice the luminosity of the $M^*_V$ ``knee'' of the normal galaxy luminosity function.  The normal galaxy luminosity function is also shown in Figure~\\ref{fig:complf}, and it is evident that it has a very different shape from that of the QSO luminosity function.  The brightest luminosity bin in the Metcalfe et al.~(1998) data extends up to $M_V \\gtrsim -24.0$, and their Schechter function is extrapolated to brighter magnitudes, as we have marked with the dashed line in Figure~\\ref{fig:complf}.  In fact, we searched for published accounts of any normal galaxies of luminosity $M_V < -24.0$ with little success.  The Sloan Digital Sky Survey has recently released a preliminary galaxy luminosity function (Blanton et al.~2001), in which the highest-luminosity bin also extends up to $M_V \\gtrsim -24.0$.  We also searched the literature on Brightest Cluster Galaxies (BCG).  In the sample of BCGs in Postman \\& Lauer~(1995), the most luminous has $M_V \\approx -23.7$, although Disney et al.~(1995) refer in passing to BCG luminosities as $M_V \\approx -24.5$.  It may be that the $M_V < -24.0$ region is dominated by QSO hosts.  We conclude that the luminosity function of QSO hosts differs greatly in shape from that of normal galaxies but that they are coincident in the highest luminosity bin, if the normal galaxy luminosity function can be extrapolated that far.  We can use the luminosity functions to estimate the ratio of the number of low-redshift QSO hosts to the number of normal galaxies as a function of absolute magnitude.  The normal galaxy luminosity function is integrated over our 0.5--mag bins and then compared to the binned host data.  The ratio of the number of hosts to the number of normal galaxies is shown in Figure~\\ref{fig:ratio}, along with a parameterization of the results.  The relationship can be expressed as $\\mathcal{R} = (L_V/{11.48L^*_V})^{2.46}$, up to $M_V \\approx -25.0$, where $\\mathcal{R}$ is the ratio of the number of hosts to the number of normal galaxies, $L_V$ is the $V$-band luminosity, and $L^*_V$ corresponds to $M^*_V=-22.35$.  The points at $M_V=-24.0$ and $-24.5$ are drawn with thin lines to indicate that they depend on the extrapolation of the normal galaxy luminosity function, although they are included in the fit.  We note that the accuracy of these results is subject to the inherent systematic uncertainties in normalization procedures for both the local galaxy luminosity function and our low-redshift QSO host galaxy luminosity function.  Normalizations of normal-galaxy luminosity functions can differ by a factor of 2 from one study to another.  The conclusions drawn here about the comparison of the QSO and normal galaxy luminosity functions are roughly similar to those reached by Smith et al.~(1986) in their ground-based study of the hosts of QSOs and lower-luminosity AGN.  The current data are of course better than what was available to them, so we believe that the conclusions are now considerably stronger.  As suggested by these authors, this ratio essentially represents an empirical parameterization of the probability that a galaxy with luminosity $L_V$ will host a QSO. It applies in the redshift interval studied here, $z\\approx 0.06-0.46$.  \\subsection {Radio-loud and Radio-quiet QSOs with Elliptical Hosts}\\label{sec:similar}  As described in \\S\\ref{sec:unweighted} and shown in Figure~\\ref{fig:MvsM} and Figure~\\ref{fig:unwLF-sub}, the QE subclass tends to have dimmer host and nuclear luminosities than the LE subclass, but there is a great deal of overlap.  Within this overlap, there must be some property, other than host luminosity, that affects the amount or nature of the fuel available to the central engine for radio emission.  The possible causes of this effect should be investigated more closely by using the members of these two subclasses of elliptical hosts to examine other properties of these QSOs, including environmental clues.  \\subsection {Selection Effects and Biases}\\label{sec:selectioneffects}  Images in our sample are of heterogeneous quality, because the original observers have chosen a variety of filters and exposure times for objects with disparate properties.  The best-exposed images, such as Q~1402$+$436 and PG~1444$+$407, average up to 15000 e pixel$^{-1}$ within the half-light radius, for a S/N up to 150 per resolution element (the full-width at half-maximum of the PSF).  The worst cases, such as 3C~93 and LBQS~0020$+$0018, have as little as 200 e pixel$^{-1}$, for a S/N of 25 per resolution element.  Adequate model fits can be obtained even in images with low S/N, as shown by the radial profile of 3C~93 (Figure~\\ref{fig:profile-3c93}). While noise does give its profile a ragged appearance, the model is not a bad fit overall, and the small discrepancy between model fit and measured profile is not reflected in the object's computed magnitude, which is based primarily on the actual counts rather than on the model fit (\\S\\ref{sec:extractmag}).  Some of the most irregularly-shaped hosts can give rise to systematic differences between the image profile and the model, as in the example of 3C~48 (Figure~\\ref{fig:profile-3c48}).  In this case, the host's shape and position, with the active nucleus offset from its center, requires masking irregular parts of the host to allow the fit to converge; this results in the systematic errors at small radii shown in the figure. But others of these irregular hosts, such as PKS~2349$-$014 (Figure~\\ref{fig:profile-pks2349}), have excellent fits, without significant systematic differences in their profiles.  Finally, Figure~\\ref{fig:profile-ms1059} shows an example of a spiral host with both the disk and bulge modeled.  A bright QSO may be expected to hide a dim host, and the more distant it is, the harder the host will be to detect.  However, we are able to see hosts that are 3.3 magnitudes dimmer than their nuclear QSO light (in apparent magnitudes in the observed wavebands).  Very few of the objects have hosts nearly this much dimmer than their nuclei, and there are in fact only 7 objects that have a host 2.1 or more magnitudes dimmer than their nucleus.  In general, considering that the host magnitudes span a noticeably smaller range than the nuclear magnitudes (Figure~\\ref{fig:MvsM}), combined with the fact that we have failed to detect a host in only one case, makes us confident that our host luminosity function is not strongly biased by missing very dim hosts.    The ellipticals outnumber the spirals in our sample, making up 62\\% of the total.  However, due to {\\it HST} target selection effects, this may not be representative of all QSOs in this redshift range.  For example, we confirm, with only one or two possible exceptions, that radio-loud QSOs are found in elliptical or interacting hosts (\\S\\ref{sec:unweighted}), while radio-quiet QSOs may be found in either ellipticals, spirals or interacting cases.  The fraction of radio-loud QSOs in our sample is $\\approx 37\\%$, higher than the 10--20\\% expected for optically--selected samples in this redshift range (Kellermann et al.~1989; Hooper et al.~1995; Hooper et al.~1996).  This may indicate that we have selected an artificially high fraction of QSOs with elliptical hosts.  But since QSOs in ellipticals are more luminous than those in spirals, our weighting procedure (\\S\\ref{sec:weightdist}) may compensate for this bias to an extent.  It is unlikely that a redshift-dependent magnitude bias could arise from non-stellar emission from the hosts.  Spiral hosts may contain H {\\sc ii} regions, however most bright H {\\sc ii} regions are masked from our fits and analyses if they are in dim areas of the host. Although less-prominent H {\\sc ii} regions might be unmasked, the rest equivalent width of H$\\alpha$, the major H {\\sc ii} emission line, is typically $\\sim 25\\,$ \\AA \\, in late-type galaxies (Gavazzi et al.~1998).  In contrast, the WFPC2 wide-band filters we use have much larger equivalent widths ranging from 867 \\AA \\, (F675W) to 1539 \\AA \\, (F814W) (Burrows~1995), so the H$\\alpha$ effect in the host is very small. More important is the affect that broad emission lines might have on a QSO's nuclear luminosity, and H$\\alpha$ is again the most prominent line.  In this case, H$\\alpha$ rest-frame equivalent widths are distributed with a median of $\\approx 240$ \\AA \\, and usually do not exceed $\\approx 450$ \\AA \\, (Sabbey~1999); thus this could account for at most a few tenths of a magnitude variation, depending on whether or not the line is included in the filter passband.  \\subsection{Comparisons with Results from Some Other Studies}\\label{sec:comparisons}  There have been a few other large-sample studies of QSO hosts with selection criteria similar to our own.  Of the space-based {\\it HST} ones, the two largest, those of Bahcall et al.~(1997) and of McLure et al.~(1999), have samples that are included in ours.  The sample of Bahcall et al.~(1997) includes 20 QSOs with redshifts $z <0.3$ and $M_V < -24.4$, making them among the most luminous objects in the nearby universe.  In this comparison, we use their final results, which come from their 2-D model fits.  Their host magnitude distribution has a shape similar to ours but is nearly a magnitude fainter.  It should be noted that the results of their 1-D fits are brighter than their 2-D fits and give a distribution about half a magnitude fainter than ours.  They find that, on average, the hosts of radio-loud QSOs are one magnitude brighter than the hosts of radio-quiet QSOs, while we find that the radio-loud hosts are about half a magnitude brighter.  They also report the hint of a luminosity difference between elliptical and spiral hosts, but they state that it may be artificial, a consequence of fitting the host model to the outer ($r\\geq 1\\farcs0$) region of the host.  Our analysis shows only a small difference, with elliptical hosts being $\\approx 0.2$ mag brighter than the spirals on average.  This is true whether using the mean or the median.  Bahcall et al.~(1997) classify 4 of their hosts as spirals, 12 as ellipticals, and the remaining 4 as interacting or of indeterminate morphology.  Ignoring the interacting and indeterminate types, we agree with their morphology classifications, with the possible exceptions of PG~1116$+$215 and PG~1444$+$407.  Bahcall et al.~(1997) classify both of these as elliptical, but we find they are disk-like. In our analysis, they are best fit by exponential profiles, and they contain central bulges, although neither host shows visual evidence of spiral arms.  One of the central conclusions of Bahcall et al.~(1997), that elliptical galaxies can host either radio-loud or radio-quiet QSOs, is confirmed by our study.  Finally, we agree with their finding that the host magnitude distribution is inconsistent with a Schechter (1976) function.  The sample of McLure et al.~(1999) includes 15 QSOs with a redshift range of $0.1 \\leq z \\leq 0.35$.  Transforming their results into $M_V$, using $\\vr=0.7$ for ellipticals and $\\vr=0.6$ for spirals (Fukugita et al.~1995), we find that their host luminosity distribution is narrower ($ -21.9 > M_V > -23.6$) than ours, though with the peak in the same magnitude bin ($M_V \\sim -23$) as ours. They find only two QSOs with spiral hosts (with the remaining 13 being ellipticals), while we have a much larger fraction of spirals, 26/70. But note that we classify MRK~1014 as a spiral on the basis of its arms, while they list it as an elliptical on the basis of its $r^{1/4}$-law radial profile.  While McLure et al.~(1999) find that essentially all radio-quiet QSOs with nuclear luminosities $M_V < -23.7$ ($M_R < -24.0$) have elliptical hosts, we find several spiral hosts with nuclei in this luminosity range.  However, our data do show that the radio-quiet QSOs with the brightest nuclei reside in elliptical hosts, while those with the dimmest nuclei reside in spiral hosts.  This may still lend support to their idea that the correlation between black hole and bulge mass derived by Magorrian et al.~(1998) affects the distribution of nuclear luminosities between elliptical and spiral hosts.  Since this investigation uses archival {\\it HST} images, we also compare our host apparent filter magnitudes, $m_{\\mathrm{host}}$, against those of the above two studies, for cases in which we use the same images.  This allows a comparison of the measurement techniques themselves.  For the Bahcall data, this corresponds to 16 objects from {\\it HST} observing programs numbered 5099, 5343, and 5849 (see Table~\\ref{table:obslist} for the object names).  For the McLure data, this corresponds to 11 objects from program number 6776 (but note that PKS 2135$-$147 is not included in their paper).  We note in particular that Bahcall et al.~(1997) provide the results of both their one-dimensional and two-dimensional model fits; while they adopt the two-dimensional magnitudes in their analysis, we find that their one-dimensional magnitudes agree better with ours, and thus adopt their one-dimensional magnitudes in the following comparison.  There are 27 other objects for which there exist published $m_{\\mathrm{host}}$ values taken from the same observations we use. These include the 9 objects observed by Boyle (program 6361) and published in Schade et al.~(2000), the 10 observed by Impey (program 5450) and published in Hooper, Impey, \\& Foltz~(1997), the four observed by Disney (program 6303) and published in Boyce et al.~(1998) [note that we exclude PG~0043$+$039, for which the published host measurements are very uncertain, and IR~0450$-$2958 and IR~0759$+$6508, which are published in Boyce et al.~(1996) without values for $m_{\\mathrm{host}}$], the two observed by Hutchings (program 5178) and published in Hutchings \\& Morris~(1995), and the two observed by Macchetto (program 5143) and published in Disney et al.~(1995) and Boyce et al.~(1998).  The comparisons show that the one-dimensional $m_{\\mathrm{host}}$ data of Bahcall et al.~(1997) are on average $0.09 \\pm 0.35$ mag fainter than ours, and the $m_{\\mathrm{host}}$ data of McLure et al.~(1999) are an average of $0.16 \\pm 0.33$ mag fainter than ours.  Looking at all 54 objects together, we find that published $m_{\\mathrm{host}}$ results average $0.30 \\pm 0.62$ magnitudes fainter than ours. Although there is considerable scatter to these differences in apparent filter magnitudes, we believe that at least part of the difference may be systematic and lie in our direct measurement of the apparent host magnitude from the PSF-subtracted image, without relying on a simple galaxy model except at large radii.  In theory, if the host model fit is weighted inversely to the square of the Poisson noise in the image ($1/\\sigma$), the model's magnitude will be slightly fainter than that of the actual host image.  In practice, we find that our host models are $\\approx 0.25$ mag fainter than the host magnitudes we calculate in \\S\\ref{sec:extractmag}.  The host magnitudes we adopt (\\S\\ref{sec:extractmag}) are not affected much by this bias in the models."
    },
    "0011/astro-ph0011549_arXiv.txt": {
        "abstract": "The tailed radio galaxies that have been called ``Type I'' are not a uniform set.  To study their dynamics, we have used the Ledlow-Owen data set, which provides a new sample of 250 radio galaxies in nearby Abell clusters.  These sources divide into two clear categories based on their radio morphology.  Type A sources (``straight'') contain nearly straight jets which are embedded in outer radio lobe. Type B sources (``tailed'') have a well-collimated jet flow which undergoes a sudden transition, at an inner hot spot, to a less collimated flow which continues on and forms a radio tail.  We have not found any separation of these classes in terms of radio power, radio flux size, galaxy power or external gas density.  We propose the difference is due to the development, or not, of a disruptive flow instability, such as Kelvin-Helmholtz, and the saturation of the instability when it develops. ",
        "introduction": "We do not yet fully understand the dynamics and evolution of radio galaxies.  Classifying sources on the basis of their morphology has been one useful approach to this question.  Such work by Fanaroff \\& Riley (``FR''; \\cite{FR}) has stimulated discussion and simple models of the sources.  FR Type II sources are generally accepted to arise from a  collimated, supersonic jet which impacts the ambient medium, resulting in a hot spot at the end of the jet, and leaving behind a cocoon which surrounds the jet.  This model is qualitatively attractive and leads to possible quantitative predictions.   Our picture of FR Type I sources, however, is much less complete.  It has been suggested that Type I's are subsonic, turbulent plumes.  This model does not, however, seem able to account for the range of phenomena seen in the data; nor has it been developed quantitatively to a point which would allow specific predictions of source evolution.  We find this frustrating, as Type I sources are much more common  in the universe than Type II's, and show much more diversity.  We present two figures which illustrate how rare classical Type II's are.  First,  one of us \\cite{JE.LF} has used the LRL, B2 and GB samples \\cite{LRL,B2,MC} to  calculate 1.4 GHz luminosity functions separately for Type II and Type I  sources, shown in Figure 1.  It is apparent that Type I's dominate Type II's below $P_{1.4} \\sim  10^{32}$W/Hz, and strongly dominate in total numbers (integrated over radio flux). Second, only 13 of 197 resolved sources in a new data set are classic II's.  This data set is the Ledlow-Owen sample of radio sources in clusters of galaxies (shown in Figure 2; \\cite{LOrad}).  Our work in this paper uses this sample.      \\begin{figure}[htb] \\begin{minipage}{0.48\\textwidth} \\resizebox{\\textwidth}{!}{\\includegraphics{Eilek1.ps}} \\caption[]{Luminosity functions for Type I and Type II sources separately.  The sources are nearby ($z < 0.5$) members of the 3CR, B2 and GB samples, identified as Type I or Type II from published radio maps.  The LF's were calculated by the inverse volume method, taking radio and (for B2) optical flux limits into account. From \\cite{JE.LF}.} \\end{minipage} \\hspace{0.04\\textwidth} \\begin{minipage}{0.48\\textwidth} \\resizebox{\\textwidth}{!}{\\includegraphics{Eilek2.ps}} \\caption[]{The full Owen-Ledlow sample in the (radio power, optical magnitude) plane, with the dividing line from \\cite{LObreak} shown.  Nearly all of the sample are in the ``Type I'' region, below the dividing line. We find that these sources should be further divided into ``Straight'' and ``Tailed'' classes.   Classic Type II's are shown as open circles.} \\end{minipage} \\end{figure}  What questions would we like to answer regarding the dynamics of radio sources?  Many spring to mind:  \\begin{itemize}  \\item How does a radio source evolve in linear size, volume, morphology?  \\item How does the luminosity of a radio source evolve with time?  \\item How can we explain the observed distribution of sources in the (size, power) or (power, galaxy magnitude) planes?  \\item What physical factors determine the morphology of a radio source?  That is, why does it become Type I or Type II in the FR classification?  Alternatively, why does it become Straight or Tailed, the types introduced in this paper?  \\item Does the morphology of a particular source change with time? Does a source change from FR type II to I, or from Straight to Tailed?  \\end{itemize}  In order to address these questions, we have undertaken a program of study based on a new data set, the Ledlow-Owen sample \\cite{LOrad} of radio sources in clusters of galaxies. We expected that the evolution of a source should be determined by such factors as its jet power, age, parent galaxy size, and the density or pressure of the surrounding medium.  To this end, we gathered radio, optical and X-ray data on the sources in the sample, and have searched for trends which might answer the questions above.  In this process, we found that the Ledlow-Owen sources fall clearly into two morphological classes. This was not what we expected. Nearly all of the sources in this sample would be called FR Type I, based on their lack of outer hot spots, and their position in the (power, magnitude) diagram (as presented by \\cite{LObreak}).  Our work with the images made it clear that the sources are not a uniform sample, but rather divide nicely into two groups.  These groups are not well described by the two well-known FR types. In fact, it now seems that the two FR types might more properly be described as  ``classic Type II's'' (the rare sources with clear outer hot spots) and ``all other sources'' (the more common sources, without outer hot spots, and including the diversity of source types we find). ",
        "conclusions": ""
    },
    "0011/astro-ph0011580_arXiv.txt": {
        "abstract": "We consider models for the generation of the emission line recently discovered in the X--ray afterglow spectrum of several bursts, and especially of GRB 991216 observed by $Chandra$. These observations suggest the presence of 0.1--1 solar masses of iron in the vicinity of the bursts. We show that there are strong geometrical and kinematical constraints on the line emitting material. We discuss several classes of models, favoring one where the line photons are produced by reflection off the walls of a wide funnel of semi--opening angle $\\theta \\approx 45^\\circ$ excavated in a young (a few months old) supernova remnant made of $\\approx 10 M_\\odot$ with radius $6\\times 10^{15}$ cm, and $\\approx 1 M_\\odot$ of iron, providing strong support for the SupraNova model. ",
        "introduction": "There are now five bursts with evidence for large amounts of material around the site of the explosion: four (GRB 970508, Piro et al., 1999; GRB 970828, Yoshida et al., 1999; GRB 991216, Piro et al., 2000, hereafter P2000; GRB 000214, Antonelli et al., 2000) show an emission feature during the afterglow, and one (GRB 990705, Amati et al., 2000) displays an edge in absorption during the burst itself. The observations of GRB 991216 by $Chandra$, begun 37~hr after the burst and lasting for about 3 hr, show a 3.49$\\pm$0.06 keV line significant at a  $\\sim 4\\sigma$ level. If the line is identified with the recombination K$\\alpha$ line from H--like iron (6.97 keV), the inferred redshift is $z= 1.00\\pm 0.02$ consistent with the largest optical absorption redshift system ($z = 1.02$) (Vreeswijk et al., 2000). The line, resolved by $Chandra$, displays a width corresponding to $0.05c$ (P2000). The line and the 2--10 keV continuum fluxes were $\\sim 1.6\\times 10^{-13}$ and $\\sim 2.4\\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$, respectively, implying (for $z=1.02$) an emission rate of iron line photons of $\\dot{N}_{\\rm line} \\simeq 4\\times 10^{52}$ s$^{-1}$, luminosity $L_{\\rm line} \\simeq 4\\times 10^{44}$ erg s$^{-1}$ and total energy $E_{\\rm line}\\sim 3 \\times 10^{49}$ erg (assuming steady emission for $40/(1+z)$~hr, with $H_0=75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$). If each iron atom produces $k$ line photons, the required iron mass is $M_{\\rm Fe} \\simeq 195\\ k^{-1}$~ M$_\\odot$. Bringing this mass down to $\\sim 0.1$~M$_\\odot$ implies $k>2000$ which in turn implies a limit on the recombination time {\\footnote{ From the formulae given in Verner \\& Ferland (1996), the recombination time can be approximated by $t_{\\rm rec}\\sim 12.7 T_7^{3/4}n_{10}^{-1}$ s in the temperature range $10^6<T<10^8$K.}}. This translates into the condition $n_{\\rm e} > 10^{10} T_7^{3/4}$ cm$^{-3}$; the electron temperature $T=10^7 T_7$ K comes from interpreting the broad bump, marginally detected by $Chandra$, at 4.4$\\pm$0.5 keV as the recombination continuum of H--like iron (9.28 keV). Analogously to GRB~950708 (Lazzati et al. 1999), the line detection implies the presence of $\\approx 0.5 M_\\odot$ of pure iron in the vicinity of the burster.  Observation of the line $t_{\\rm obs}$ after the burst implies that the line emitting material must be within $\\sim ct_{\\rm obs}/(1+z)$ from the burst site: this region must then be {\\it compact}, contain $\\sim 0.1 M_\\odot$ only in iron, but nevertheless be optically thin to electron scattering, such that Comptonization does not broaden the line beyond the observed width. We call this the {\\it size problem}.  Furthermore, if we interpret the line width as due to the velocity of a supernova remnant, the limit on the size allows to estimate the age of the remnant: for GRB 991216,  $\\sim 15$~days. At this time (Fig. 1), cobalt nuclei outnumber nickel and iron nuclei and the line would be produced mainly by cobalt, not iron, at an energy $\\epsilon = 7.5/(1+z)$ keV. We call this the {\\it kinematic problem}.  We discuss in section 2 the constraints implied by the line observation; we then present in section 3 three scenarios which could produce the observed line. Two of them are found to be viable only by adopting ad hoc assumptions or through fine tuning, while a third scenario, the ``wide funnel\" scenario, appears to be more promising. In section 4 we discuss our results. ",
        "conclusions": "X--ray line emission features from GRBs' afterglows give us information on the nature of their progenitors by imposing constraints on models, the most severe being how to arrange a large amount of line emitting material around the GRB site, but without large Thomson scattering opacities. A further limit comes from the line width of $Chandra$'s observation in GRB 991216. These observations require a large iron mass (Eq. \\ref{rate}), such as is seen only in SNe.  We have presented three distinct models: two (the ``back of the remnant\" and ``back illuminated equatorial material\" models) require that the observed emission line is produced by redshifted cobalt or nickel (instead of iron), and need fine tuning and/or ad hoc assumptions. The ``wide funnel\" model, instead, solves the size problem, and the acceleration of the line emitting material by grazing incident photons solves the kinematic problem, allowing time for cobalt to decay to iron. It implies that the GRB progenitors are massive stars exploded as supernovae months before the burst, contaminating the circum--burst environment with iron rich material. This two--step process, and the time--delay between the two steps, are predicted by the SupraNova scenario (Vietri \\& Stella, 1998). Alternative explanations invoke extremely iron enriched massive winds (Weth et al. 2000).  Future observations may address issues such as the presence of cobalt and nickel lines in the spectrum, the age of the remnant, the line profile, the geometry and kinematics of the emitting region, the presence of more than one line, the ionization state, the time evolution of the line and edges, all indicative of the characteristics of the illuminator. Together, this information may determine whether the line emitting material originates in a normal, albeit rare, supernova remnant or if a more exotic explosive phenomenon is required."
    },
    "0011/astro-ph0011313_arXiv.txt": {
        "abstract": "We review the present status of the Baikal Neutrino Project and present preliminary results of a search for upward going atmospheric neutrinos, WIMPs and magnetic monopoles obtained with the  detector {\\it NT-200} during 1998. Also the results of a search for very high energy neutrinos with partially completed detector in 1996 are presented.  \\vspace{1pc} ",
        "introduction": "The Baikal Neutrino Telescope  is deployed in Lake Baikal, Siberia, \\mbox{3.6 km} from shore at a depth of \\mbox{1.1 km}. {\\it NT-200}, the medium-term goal of the collaboration \\cite{APP}, was put into operation at April 6th, 1998 and consists of 192 optical modules (OMs). An umbrella-like frame carries  8 strings, each with 24 pairwise arranged OMs. Three underwater electrical cables and one optical cable connect the detector with the shore station.  The OMs are grouped in pairs along the strings. They contain 37-cm diameter {\\it QUASAR} - photo multipliers (PMs) which have been developed specially for our project \\cite{OM2,Project}. The two PMs of a pair are switched in coincidence in order to suppress background from bioluminescence and PM noise. A pair defines a {\\it channel}.  A {\\it muon-trigger} is formed by the requirement of \\mbox{$\\geq N$ {\\it hits}} (with {\\it hit} referring to a channel) within \\mbox{500 ns}. $N$ is typically set to \\mbox{3 or 4.} For  such  events, amplitude and time of all fired channels are digitized and sent to shore. A separate {\\em monopole trigger} system searches for clusters of sequential hits in individual channels which are characteristic for the passage of slowly moving, bright objects like GUT monopoles.  Here we present preliminary results of analysis of data, which were accumulated in the first 234 live days of {\\it NT-200} as well as results obtained from the analysis of data taken with {\\it NT-96}, the 1996 stage of the detector. ",
        "conclusions": "The deep underwater neutrino telescope {\\it NT-200} in Lake Baikal is taking data since April 1998. Using the first 234 live days, 35 neutrino induced upward muons have been reconstructed. Although in a good agreement with MC expectation this number is on factor 3 lower then predicted for the fully operational {\\it NT-200}. The reason is that,  due to unstable operation of electronics, in average only 50 - 70 channels have taken data during 1998. This is in contrast to 1999 and 2000 data taking where stability had improved. Ten events within a 30 degree half angle cone around nadir have been selected and limits on the excess of muon flux due to WIMP annihilation in the center of the Earth have been derived. Also a new limit on the flux of fast monopoles has been obtained.  In the following years, {\\it NT-200} will be operated as a neutrino telescope with an effective area between 1000 and 5000 m$^2$, depending on the energy. It will investigate atmospheric neutrino spectra above 10 GeV (about 1 atmospheric neutrino per two-three days). Due to the high water transparency and low light scattering with effective scattering length greater than 150m$\\div$200m, the effective volume of {\\it NT-200} for high energy electron and tau neutrinos detection is more than two orders of magnitude larger than its geometrical volume. This will permit a search for diffuse neutrino fluxes from AGN and other extraterrestrial sources on a level of theoretical predictions.  With an effective area two times larger than Super-Kamiokande, for nearly vertically upward muons (E$_{\\mu}>$10 GeV) {\\it NT-200} will be one of the most powerful arrays for indirect search for WIMP annihilation in the center of the Earth during the next few years. It will also be a unique environmental laboratory to study water processes in Lake Baikal.  Apart from its own goals, {\\it NT-200} is regarded to be a prototype for the development a telescope of next generation with an effective area of 50,000 to 100,000 m$^2$. The basic design of such a detector is under discussion at present.  \\bigskip  {\\it This work was supported by the Russian Ministry of Research,the German Ministry of Education and Research and the Russian Fund of Fundamental Research ( grants } \\mbox{\\sf 99-02-18373a}, \\mbox{\\sf 00-02-31001} {\\it and} \\mbox{\\sf 00-15-96794}) {\\it and by the Russian Federal Program ``Integration'' (project no.} 346)."
    },
    "0011/astro-ph0011125_arXiv.txt": {
        "abstract": "We present an analysis of high quality spectra of members of the young cluster M35. By using a multi-fiber spectrograph, we are able to collect high signal--to--noise, high resolution  spectra of a sample of photometric candidate members. Accurate radial velocities are used to establish the membership status and rotational velocities are measured using cross-correlation. We also derive the metal content of the cluster,   [Fe/H]$_{\\rm M35}$=--0.21$\\pm$0.10, based on spectral synthesis. Finally, we  derive the lithium abundances of the {\\it bona fide} cluster members and compare the results with members of other clusters. For example, M35 shows a smaller range in both rotation rates and lithium abundances as compared to the Pleiades.  We discuss possible roles of various parameters. Our high quality M35 database of lithium abundances and rotational velocities are perfectly suited to be used as a laboratory to test theoretical models dealing with the lithium depletion phenomenon. We discuss the role of stellar inhomogeneities and rotation on the lithium depletion phenomenon. ",
        "introduction": "Late spectral type  stars share a series of characteristics which make them extremely interesting objects. Among other parameters, these properties usually depend on the age and the stellar mass. These stars have activity which is analog to the one present on the Sun, arising from the stellar corona, the chromosphere and the  photosphere.  This activity is produced as a consequence of the interaction  between differential rotation and magnetic fields, the so-called stellar dynamo (Parker 1955). Rotation itself is a very important characteristic, with a dependence on mass and age not fully understood. The understanding of these properties require to locate them in an evolutionary scenario. This is done by studying members of open clusters of different ages and metallicities, so we can determine which is the particular role of age, stellar mass and chemical composition.  Lithium is a chemical element which can be used as a probe to the internal structure, since convection transports the material downward to layers hot enough  to destroy lithium by interactions with protons. Due to this, its surface abundance depends strongly on age for a given mass during the pre-Main Sequence (pre-MS) phase. And, because of increasing depth of the convective envelope when moving to less massive stars, the lithium abundance also depends on mass. (or stars of spectral type later than F8, at a given age, more massive stars have higher Li.) After the pre-MS phase, additional mixing mechanisms are needed, in order to explain the lithium abundance of FGK spectral type members of clusters of different ages. For recent reviews, see Deliyannis (2000) and Jeffries (2000). In particular, Deliyannis (2000) stresses  the stellar mass as the first parameter controlling Li depletion, and age the second parameter.  Additional parameters might be composition, and initial angular momentum. There have been reported a significant amount of studies on the evolution of lithium with age and the dependence on mass. Updated reviews on the lithium phenomenon, from the theoretical and observational point of views,  can be found in Vauclair (2000) and Mart\\'{\\i}n (2000), respectively. Very well known open clusters such as the Hyades,  Praesepe, Coma, M67 and Alpha Per  have been exhaustively investigated, as well as other not so famous, and the behavior of lithium in their members analyzed. The Pleiades has played a {\\it prima donna} role in these investigations. Duncan \\& Jones (1983) showed that G Pleiades stars have significant lithium abundances differences. Butler et al. (1987) reported a scatter for rapid rotators of  K spectral type, confirmed by Garc\\'{\\i}a L\\'opez et al. (1991a,b),  and Soderblom et al. (1993a).  Garc\\'{\\i}a L\\'opez, Rebolo, \\& Mart\\'{\\i}n (1994) added more data to the sample and established that the relation between fast rotators and high lithium abundances broke down for stars at Teff$<$4500 K. Jones et al. (1996) examined the late K and Russell (1996), Stuik et al. (1997), Jeffries (1999), King, Krishnapurthi \\& Pinsonneault (2000) and Barrado y Navascu\\'es et al. (2000) have recently  dealt with the origin of the lithium spread in K Pleiades stars, without solving completely the puzzle.   The M35 cluster (NGC2168) is, allegedly,  coeval to the Pleiades, since their turn--off is located in the same position in the color-magnitude diagram (Vidal 1973). Due to several properties, it  has  a promising future as a target for this type of studies. Although it is moderately far away --(m-M)$_0$=9.7 magnitudes, compared with (m-M)$_0$=5.36 in the case of the Pleiades, Robichon et al. (1999)--, it has a clear advantage over other associations: its richness. It is one of the richest nearby young clusters, with a total mass estimated between 1600 and 3200 M$_\\odot$ (Leonard \\& Merritt 1989). Membership of those (few) stars hotter than about 5700 K can be checked against the excellent proper motion membership study of McNamara and Sekigushi (1985). In addition, several optical  surveys have been recently conducted, producing long lists of candidate members (Sung \\& Bessell 1999; von Hippel 2000; Sarrazine et al. 2000). The recent study by Sung \\& Bessell (1999) yields (m-M)$_0$=9.60$\\pm$0.10, E(B--V)=0.255$\\pm$0.024 and an age of 200$_{-100}^{+200}$ Myr, somewhat older than the turn--off age of the Pleiades, 80 Myr. Similar results have been derived by Sarrazine et al. (2000), including (m-M)=10.16$\\pm$0.01, E(B--V)=0.198$\\pm$0.008, and an age of 160$\\pm$40 Myr.   For some time now, we have been studying exhaustively this cluster, tackling problems such as the mass function of the cluster (Barrado y Navascu\\'es et al. 1999, 2000a), rotational periods, and  multicolor photometric surveys (Barrado y Navascu\\'es et al. 2000a). Moreover, since the cluster covers less than 1 sq.deg., it is a perfect candidate for multi--object spectroscopy. We have taken advantage of the WIYN/HYDRA multi--fiber spectrograph and observed a sample of M35 candidate members, with the goal of identifying {\\it bona fide} members, determining the pattern of lithium depletion and rotation for them, and establishing  the connection between these two last properties (rotation and lithium).  We present the data in Section 2, whereas the analysis and comparison with open clusters, including the Pleiades, are discussed in Section 3. A summary and the main  conclusion can be found in Section 4. ",
        "conclusions": "We have collected multi-fiber, high resolution  spectroscopy of 76 candidate members of the young cluster M35. Our spectra have moderate to high signal--to--noise rations (S/N$\\sim$40--160), a remarkable feat for stars located at (m-M)=10.4. We have measured radial velocities and using these values as membership criterion, we have cataloged 39 stars out of the original 76 as probable members of the cluster. Another 13 stars are spectroscopic binaries and part of them could be members of the cluster. We also have measured rotational velocities via cross--correlation. The comparison between the the rotation pattern of members of M35 and the Pleiades indicates that, indeed, M35 is somewhat older, about 175 Myr, since it lacks fast rotators. Our high quality M35 spectra have been used to derive the metallicity of the brighter and warmer members of the cluster. The result yields  [Fe/H]$_{\\rm M35}$=--0.21$\\pm$0.10, lower than the Pleiades metal content, [Fe/H]=+0.01. Subsequently, we have derived  lithium abundances of the M35 members, and compared these values with members of other nearby clusters of different ages and metallicities. We conclude that  at the age of M35, most of the lithium spread observed in Pleiades stars is not longer present. We have interpreted this fact, at least partially,  as an effect of the photospheric spots on the apparent lithium abundance of Pleiades stars. Stellar rotation could also play some role. The data presented in this study are perfectly suited to be used as a validity check of theoretical models involving the lithium depletion."
    },
    "0011/astro-ph0011369_arXiv.txt": {
        "abstract": "We present preliminary results of a follow-up survey which aims to characterise in detail those galaxies which hosted Type Ia supernovae found by the Supernova Cosmology Project. Our survey has two components: Hubble Space Telescope imaging with STIS and Keck spectroscopy with ESI, the goal being to classify each host galaxy into one of three broad morphological/spectral classes and hence to investigate the dependence of supernovae properties on host galaxy type over a large range in redshift. Of particular interest is the supernova Hubble diagram characterised by host galaxy class which suggests that most of the scatter arises from those occurring in late-type irregulars. Supernovae hosted by (presumed dust-free) E/S0 galaxies closely follow the adopted SCP cosmological model. Although larger datasets are required, we cannot yet find any significant difference in the light curves of distant supernovae hosted in different galaxy types. ",
        "introduction": "The discovery and controlled systematic study of faint, distant type Ia supernovae (SNe\\,Ia) has transformed the prospects for constraining the time derivative of the cosmic expansion rate.  Though well-measured local SNe\\,Ia have a `raw' observed dispersion in their maximum luminosity of $\\simeq0.35\\,\\rm{mags}$, empirical correlations between the peak SNe luminosity and light-curve width enable this scatter to be reduced (Phillips 1993; Hamuy et al. 1996b; Riess et al. 1996; Perlmutter et al. 1999). By using such light curve shape corrections with independent samples of SNe\\,Ia, two Hubble diagrams have been constructed by the High-Redshift Supernovae Search Team (Riess et al. 1998) and the Supernovae Cosmology Project (SCP, Perlmutter et al. 1999, hereafter P99). These data strongly exclude the hitherto popular Einstein de Sitter cosmology ($\\Omega=1$, $\\Lambda=0$). In combination with the results of recent microwave background measurements (de Bernardis et al. 2000; Jaffe et al. 2000) which indicate a spatially flat inflationary universe, the SNe\\,Ia results suggest a significant non-zero cosmological constant ($\\Omega=0.28$, $\\Lambda=0.72$).  Conclusions of this importance require excellent supporting evidence. In particular, it is appropriate to question the homogeneity, environmental trends and evolutionary behaviour of the SNe found at all redshifts. Systematic differences between high and low redshift samples might change the derived cosmological parameters without destroying the small dispersion seen in the SNe\\,Ia Hubble diagrams. Evolutionary differences between low and high-redshift SNe\\,Ia might arise via the progenitor composition (H\\\"{o}flich 1999), a differential dust extinction with greater amounts of dust in high-redshift environments, either in the host galaxy (Totani \\& Kobayashi 1999) or `grey' dust in the IGM (Aguirre 1999), or a dependence of the SNe properties on host galaxy environments. Addressing this last possibility forms the basis of this present study.  Though SNe\\,Ia can occur in all types of galaxies, disk and spheroidal stellar populations should sample different ranges of metallicities and dust content, and thus we might expect that SNe\\,Ia progenitor composition and the light curve properties could be affected accordingly. By studying SNe grouped according to the mean underlying stellar population, we can directly investigate such effects.  Local studies of the environmental effects of local SNe\\,Ia have been conducted by Hamuy et al. (1996a, 2000), Branch et al. (1996) and Riess et al. (1999). In various low-redshift samples, they find the distribution of light curve decline rates correlates with the morphological type of the host galaxy, with ellipticals hosting faster decline rate (dimmer) SNe, whereas spiral galaxies preferentially host slower decline rate (brighter) SNe.  Riess et al.  (1999) also find that SNe\\,Ia with fast decline rates occur further from the host galaxy centre.  Hamuy et al. (2000) present results from a larger combined SNe sample, and find that brighter SNe occur in bluer stellar environments (also noted by Branch et al.  1996), as well as some evidence to suggest that the brightest SNe occur in the least luminous galaxies. We note that no claim has yet been made that such correlations would apply in such a manner as to reduce the significance of a non-zero cosmological constant (Hamuy et al. 2000).  Here, we extend these studies to include similar comparisons between the host galaxies and properties of high-redshift SNe, where environmental effects on the determination of the cosmological parameters can be more easily explored. We present the preliminary results of our study probing the environments of high-redshift SNe\\,Ia discovered via the SCP, and compare these high-redshift SNe/host galaxy correlations with the local sample of SNe of Hamuy et al.  (1996a), the same sample previously used in P99 for the determination of cosmological parameters. Though larger local SNe\\,Ia samples have recently become available (Riess et al.  1999), in this preliminary analysis we restrict ourselves to the above local sample for ease of comparison with the P99 cosmological results, deferring a full analysis to a later article (Sullivan et al., in prep.). ",
        "conclusions": "The preliminary results of our investigation continue to support the conclusions of P99. However, the Hubble diagram categorised by galaxy class offers new insight into the origin of the scatter in the diagram. We find the scatter about the best-fit cosmology is greater when determined from SNe which occurred in late-type/irregular galaxies. Minimal scatter is found for SNe drawn from E/S0 galaxies as expected if dust is largely absent in these systems. A non-zero $\\Lambda$ is supported by SNe arising from each galaxy type individually, with late-type galaxies implying a larger $\\Lambda$ than E/S0 galaxies as might be expected when simple dust models are considered. Finally, little difference is found when the Hubble diagram is categorised according to SNe located at large and small projected distances from their host galaxy.  \\begin{figure} \\centering \\epsfig{file=iau_ellis_5.eps,width=3in} \\caption{ The distribution of SNe stretch with projected distance from host galaxy centre. Only galaxies with a \\textit{HST} image are shown. The boxed points refer to SNe excluded from the primary fit of P99 either because they are stretch or residual outliers. We see a slight trend with SNe located at a greater projected distance having a dimmer peak luminosity, similar to fig.~5c in Riess et al. (1999).} \\label{fig_stretch_distance} \\end{figure}  In the future, we will attempt to further subdivide our emission line galaxies into two further types according to the spectral characteristics. By using the high resolution of the ESI, we can classify emission line galaxies according to the spatial variance of their H$\\alpha$/H$\\beta$ ratio, describing how the dust extinction in such objects varies according to the location inside the galaxy. We will attempt to generate a subsample of these objects which posses a dust distribution that shows large variances and examine whether these objects are responsible for the significant scatter in the current Hubble plot (Figure~3) among the late-type emission line galaxies. Those with smoother dust distributions might be expected to have less intrinsic scatter in the SNe peak luminosities (due to uncertain dust corrections), typical of those SNe arising from E/S0 galaxies.  While these initial results are promising, the current sample size is only around half that used in the full Hubble diagram in P99, and it is clear that further \\textit{HST} imaging and spectroscopic studies of the host galaxy population are needed to confirm the initial findings presented here."
    },
    "0011/astro-ph0011475_arXiv.txt": {
        "abstract": "We discuss how measurements of the absorption of $\\gamma$-rays from GeV to TeV energies via pair production on the extragalactic background light (EBL) can probe important issues in galaxy formation. We use semi-analytic models (SAMs) of galaxy formation, set within the hierarchical structure formation scenario, to obtain predictions of the EBL from 0.1 to 1000$\\mu$m. SAMs incorporate simplified physical treatments of the key processes of galaxy formation --- including gravitational collapse and merging of dark matter halos, gas cooling and dissipation, star formation, supernova feedback and metal production --- and have been shown to reproduce key observations at low and high redshift.  Here we also introduce improved modelling of the spectral energy distributions in the mid-to-far-IR arising from emission by dust grains.  Assuming a flat \\lcdm\\ cosmology with $\\Omega_m=0.3$ and Hubble parameter $h=0.65$, we investigate the consequences of variations in input assumptions such as the stellar initial mass function (IMF) and the efficiency of converting cold gas into stars. We conclude that observational studies of the absorption of $\\gamma$-rays with energies from $\\sim$10 Gev to $\\sim$10 TeV will help to determine the EBL, and also help to explain its origin by constraining some of the most uncertain features of galaxy formation theory, including the IMF, the history of star formation, and the reprocessing of light by dust. ",
        "introduction": "The extragalactic background light (EBL) represents all the light that has been emitted by galaxies over the entire history of the universe. The EBL that we observe today is an admixture of light from different epochs, its spectral energy distribution (SED) distorted by the redshifting of photons as they travel to us from sources at different distances. It is therefore a constraint on both the intrinsic SEDs of the sources and their distribution in redshift.  At present, there is more than a factor of two uncertainty in the amplitude of the EBL in the UV, optical, and near-infrared \\cite{puget}. The EBL in the mid-IR is even more uncertain. The far-IR background measured at $\\gsim 100 \\mu$m \\cite{puget96,guider97,hauser,fixsen} represents at least half of the total energy in the EBL, yet the sources that produced it remain uncertain.  High energy $\\gamma$-ray astronomy promises to help resolve these uncertainties by providing independent constraints on the EBL, in the mid-IR with $E_\\gamma$ in the $\\sim 10$ TeV energy range, and in the 0.1-3 $\\mu$m range with $E_\\gamma \\sim100$ GeV via the new low-threshold instruments that will soon be available.  High energy $\\gamma$-rays from sources at cosmological distances are absorbed via electron-positron pair production on the diffuse background of photons that comprises the EBL. Thus, $\\gamma$-ray observations of objects with known redshift and intrinsic spectral shape will constrain the EBL in these crucial wavelength regimes by measuring the optical depth of the Universe to photons of various energies.  This in turn will help to constrain some of the most fundamental uncertainties in physical models of galaxy formation.  In order to illustrate this, in this paper we use a ``forward evolution'' approach, which attempts to model the essential features of galaxy formation using simple recipes. These semi-analytic models are set within the modern Cold Dark Matter (CDM) paradigm of hierarchical structure formation, and trace the gravitational collapse and merging of dark matter halos, the cooling and shock heating of gas, star formation, supernovae feedback, metal production, the evolution of stellar populations and the absorption and re-emission of starlight by dust. This machinery has been used extensively to predict optical properties of low-redshift galaxies, with good results (e.g., \\cite{kwg,cole94}; reviewed and extended in \\cite{sp,spf}, hereafter SP and SPF).  A semi-analytic approach was also used by Devriendt and Guiderdoni \\cite{dev2} to make predictions of counts and backgrounds in the mid-to-far-IR, with more detailed modelling of dust extinction and emission, but less detailed modelling of merging and star formation.  We have now combined the strengths of these two models, by integrating the stellar SEDs and dust modelling of \\cite{dev1,dev2} into the galaxy formation code of the Santa Cruz group.  Some parts of the ``standard paradigm'' of galaxy formation represented by our SAMs are relatively solid. For example, once a cosmological model and power spectrum are specified, it is straightforward to compute the gravitational collapse of dark matter into bound halos using $N$-body techniques, and analytic formalisms such as those used in our modelling \\cite{sk} have been checked against these results \\cite{slkd}. Within the range of values for the cosmological parameters allowed by existing observational constraints (i.e., $\\Omega_{\\rm matter} \\simeq 0.3-0.5$, $\\Omega_{\\rm matter}+\\Omega_\\Lambda \\simeq 1$, $H_0 \\simeq 60-80$ km/s/Mpc; see e.g. \\cite{primack2000} for a summary), these results do not change significantly.  Similarly, modelling of gas cooling appears to be fairly robust and agrees well with hydrodynamic simulations \\cite{pearce}. However, other aspects, notably the efficiency of conversion of cold gas into stars, the effect of subsequent feedback due to supernovae winds or ionizing photons, the stellar initial mass function (IMF), and the effects of dust, remain highly uncertain, and some predictions are quite sensitive to their details.  For example, SPF showed that the star formation history of the Universe and the number density of high redshift $z \\gsim 2$ ``Lyman-break'' galaxies (LBGs; e.g. \\cite{steidel:99}) may be quite different depending on whether star formation is primarily regulated by internal properties, such as gas surface density in a quiescent disk, or triggered by an external event such as an interaction. Because the largest samples of LBGs are primarily identified in the rest UV, model predictions are also quite sensitive to the high-stellar-mass slope of the IMF, and to dust extinction. At the other end of the spectrum is the sub-mm population detected by SCUBA, believed to be predominantly high redshift ($z \\gsim 2$) luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs) powered by star formation rates of hundreds to thousands of solar masses per year (e.g., \\cite{sanders}). Theoretical predictions of the numbers and nature of these objects are highly sensitive to the same issues (the dominant mode of star formation, dust, the IMF), but provide a crucial counter-balance to the optical observations. However, the current mismatch between the sensitivity and spatial resolution of optical and sub-mm instrumentation has made it difficult to establish the connection between the two populations observationally.  The Milky Way, like most nearby galaxies, emits the majority of its light in optical and near-IR wavelengths; only about 30\\% of the bolometric luminosity locally is released in the far-infrared \\cite{sn:91}.  This was generally believed to be typical of most of the starlight at all redshifts until the discovery of the far-IR part of the EBL by the DIRBE and FIRAS instruments on the COBE satellite, at a level ten times higher than the no-evolution predictions based on the local luminosity function of IRAS galaxies, and representing twice as much energy as the optical background obtained from counts of resolved galaxies \\cite{madaupoz}. This result suggests that either the dust extinction properties of ``normal'' galaxies change dramatically with redshift, or a population of heavily extinguished galaxies (perhaps analogous to local LIRGs and ULIRGs) is much more common at high redshift than locally, or both.  Some of these galaxies may have already been observed, at 15 $\\mu$m by ISO \\cite{elbaz99}, and at 850 $\\mu$m by SCUBA \\cite{blain}.  Guiderdoni et al. \\cite{ghbm,dev2} showed that their simplified semi-analytic model could reproduce the multi-wavelength data only if they introduced a population of heavily extinguished galaxies with high star formation rates, and with strong evolution of number density with redshift.  This population was introduced ad-hoc by \\cite{ghbm,dev2}, but as discussed by these authors, by \\cite{silkdev} (based on \\cite{bss}), and also by SPF, the increasing importance of starbursts at high redshift, due to the increasing merger rate and higher gas fractions, is a natural mechanism to produce this population. The models of SPF contain a detailed treatment of mergers and the ensuing collisional starbursts, which has been calibrated against the merger rate in cosmological $N$-body simulations \\cite{kolatt} and the starburst efficiency in hydrodynamical simulations \\cite{mh,somerville}.  Moreover, they produced good agreement with observations of LBGs and Damped Lyman-$\\alpha$ systems (SPF) as well as low redshift galaxies (SP). Therefore, it will be extremely interesting to see if these same models, when combined with the more sophisticated treatment of dust extinction and emission developed by Devriendt, Guiderdoni, and collaborators, will be able to simultaneously reproduce observations over the broad range of wavelengths and redshifts discussed above.  In the next section we briefly describe the ingredients of our models, and then present the results of the predicted EBL.  The following section presents the implications for $\\gamma$-ray attenuation, and the final one briefly discusses some alternative treatments and our own conclusions.  The work summarized here is a brief, preliminary sample of the results which will soon be presented in a series of papers, now in preparation, on the EBL and its breakdown into various kinds of sources and on the implications for $\\gamma$-ray astronomy. ",
        "conclusions": ""
    },
    "0011/astro-ph0011533_arXiv.txt": {
        "abstract": "s{There is a tension between observed Helium abundance and the prediction of the standard Big Bang Nucleosynthesis. We show that non-minimally quintessence model may help to reduce this tension between theory and observation.}  Recently, it has been discovered that the expansion of the Universe is accelerating\\cite{SNIa}. This requires the existence of a dark energy component in the Universe with an equation of state $p = w \\rho$, $w<0$. One example of such a component is a cosmological constant. However, it is difficult to understand in the present framework of particle physics why it is so small. A fine tuning of $10^{-120}$ is required if the cosmological constant arises from Planck scale physics.  Quintessence models\\cite{quintessence} were suggested as an alternative to cosmological constant. In quintessence models, a scalar field provides the dark energy which drives the accelerated expansion of the Universe. The evolution of the scalar field is such that its equation of state mimics the dominant component of the Universe, thus explains why it is so small at present time.  In Non-minimally Coupled (NMC) quintessence models\\cite{NMC} the scalar field couples to the gravitational constant. The action of NMC can be written as \\begin{equation} S=\\int d^{4} x \\sqrt{-g} \\left[\\frac{1}{2}F R-\\frac{1}{2}\\phi^{;\\mu} \\phi_{;\\mu} -V(\\phi) + L_{\\rm fluid}\\right],\\nonumber \\label{NMCeq} \\end{equation} where \\begin{equation} F(\\phi)=1-\\xi(\\phi^2 - \\phi_0^{2}),\\quad V(\\phi)=V_{0} \\phi^{-\\alpha}. \\end{equation} Such coupling could arise if for example, the scalar field is a dilaton in superstring theory. One of the motivations for introducing such non-minimal coupling is to address the ``coincidence problem'': why does the dark energy component happens to become dominant at the present epoch? Had this occurred at an earlier epoch, growth of structure due to gravitational instability would be inhibited, and any life form would be impossible to exist. By introducing a coupling between curvature and the quintessence field, it was hoped that the scalar field dominance could be triggered automatically shortly after the Universe becomes matter dominant. Unfortunately, for the NMC models discussed here, it was found that the trigger mechanism does not work\\cite{NMC,Liddle-Scherrer}, nonetheless, the coupling to gravity could have other interesting consequences. Here, I show that NMC models provides a possible solution of the ``helium problem'' in the big bang nucleosynthesis.   The standard model of Big Bang Nucleosynthesis (BBN) is an enormously successful theory. The predicted abundances of the light elements, which ranges ten orders of magnitude, were found to be consistent with observations\\cite{OSW00}. In particular, the BBN prediction of $~^{4}\\mathrm{He}$ abundance ($Y_p \\approx 0.25$) provides the first evidence of a hot big bang beginning of the Universe.  The $~^{4}\\mathrm{He}$ abundance were mostly influenced by two factors, the expansion rate of the Universe during BBN, and the baryon to photon ratio $\\eta$  (see Fig.~\\ref{Y-eta}). The helium abundance increases with the expansion rate for two reasons. First, BBN starts when the weak interaction which converts proton to neutron ceases to be effective. This occurs when $\\Gamma \\sim H$, where $\\Gamma, H$ are the reaction rate and expansion rate, respectively. For a higher $H$ at a given scale factor $a=1/(1+z)$, BBN starts earlier, when the neutron fraction is higher. Second, this also meant a shorter interval for the neutrons to decay before it is combined in subsequent nuclear fusion. Both of these two effects enhance the neutron fraction. Since most of these neutrons ended up in $~^{4}\\mathrm{He}$, a faster expanding universe would yield more helium. Thus, once $\\eta$ is determined from either deuterium abundance or other methods such as cosmic microwave background (CMB) anisotropy, the $~^{4}\\mathrm{He}$ abundance could be used to constrain the expansion rate during BBN.  The helium abundance in extragalactic HII (ionized hydrogen) regions could be obtained by observation of the HeII $\\to$ HeI recombination lines. Since $~^{4}\\mathrm{He}$ is also produced in stars along with heavy elements such as Oxygen, it is expected that the primordial $~^{4}\\mathrm{He}$ abundance could be obtained by extrapolation to zero Oxygen abundance. Using this technique, Oliver and Steigman obtained \\cite{Oliver-Steigman} \\begin{equation} Y_p = 0.234 \\pm 0.003 (stat.), \\end{equation} while Izotov and Thuan obtained a higher value\\cite{Izotov-Thuan} \\begin{equation} Y_p = 0.244 \\pm 0.002 (stat.). \\end{equation} Clearly these two data sets are statistically inconsistent with each, due to large systematic errors. Below, we adopt a midway value of \\begin{equation} Y_p = 0.239 \\pm 0.005, \\label{midway} \\end{equation} or, $0.229< Y_p < 0.249$ at 95\\% C.L.  \\begin{figure}[t] \\epsfxsize=4in % \\epsfbox{Y-eta.ps} % \\caption{Helium abundance vs. $\\eta$. The solid curve is the prediction of standard BBN with 3 neutrinos, the two short-dashed curves are standard BBN with 2 and 4 neutrinos, and the long-dashed curve is the result for a NMC model. The horizontal lines mark the center and 2$\\sigma$ limits of helium in our ``midway'' approach. The vertical lines mark the value of $\\eta$ determined from combined COBE+Boomerang+Maxima data.\\label{Y-eta}} \\end{figure}   In order to compare theory with observation, we also need to determine $\\eta$. $\\eta$ could be determined from BBN. Burles and Tytler\\cite{Burles-Tytler} obtained D/H=$(3.3 \\pm 0.25) \\times 10^{-5}$, corresponding to a lower bound on $\\eta$ at $2\\sigma$ level \\begin{equation} \\eta_{10} \\equiv 10^{10} \\eta <6.3. \\end{equation} CMB anisotropy provides another way of measuring baryon density. Recently, an analysis of the combined data from Boomerang and Maxima yields a higher baryon density. The best fit model with a flat universe yields\\cite{Jaffe} \\begin{equation} \\Omega_b h^2 =0.030 \\pm 0.004, \\qquad \\eta_{10} = 8.2\\pm 1.0, \\end{equation}  If we assume that there are three neutrino species, and adopt $\\eta \\approx 4.5$ as inferred from the deuterium abundance, then the standard BBN $~^{4}\\mathrm{He}$ abundance is in disagreement with the results of Oliver and Steigman. It is in marginal agreement with the ``midway'' result of Eq.~\\ref{midway}, but still at high the end. If we adopt the $\\eta$ inferred from CMB, then even the ``midway'' limit is exceed (see Fig.~\\ref{Y-eta}).  Furthermore, in addition to the three standard model neutrinos, a sterile neutrino may be needed to explain the results from neutrino oscillation experiments\\cite{neutrino}. If either or both of these were confirmed, or if there is any other light particle in the Universe, the breach between theory and observation on $~^{4}\\mathrm{He}$ would become even wider.  How could we make a model which produce less helium? If the expansion of the Universe is slower at the time of BBN, then the helium abundance is reduced. In the standard BBN model, the expansion rate is given by \\begin{equation} H^2 = \\frac{8\\pi G}{3} \\rho, \\end{equation} Thus, $\\rho$ becomes greater with the introduction of each new particle species.  In quintessence models, $\\rho_{tot} = \\rho_f + \\frac{1}{2}\\dot{\\phi}^2 + V(\\phi)$, is it possible to introduce a negative $V$ to reduce $\\rho$? Unfortunately, this would not work. To see this, note that if a negative potential is introduced, the minima of the potential must have $V<0$, the Universe would fall to this potential well. Since $\\rho_f$ is a decreasing function of $a$, sooner or later we would reach to the point that $\\rho_{tot} = 0$, further expansion is not possible, and the Universe would begin to contract. Such a contraction of the Universe in the future is not ruled out, however, if we hope to use a negative potential to reduce helium production, the negative potential must become sub-dominant at BBN era, and then become dominant well before the current era, which is incompatible with observation.  With the NMC model introduced in Eq.~\\ref{NMCeq}, however, it is possible to obtain a lower helium abundance, because now we have \\begin{equation} H^2 = \\frac{8\\pi G}{3F} \\left(\\rho+\\frac{1}{2}\\dot{\\phi}^2 + V(\\phi) -3 \\dot{F} H \\right). \\end{equation} For $F>1$ the value of $H$ could be lowered.  \\begin{figure}[t] \\epsfxsize=4in % \\epsfbox{alphazeta-xif.ps} % \\caption{$\\zeta-1 $ as a function of $\\alpha$, the five curves from top to bottom are for models with $\\xi=0.004, 0.008, 0.012, 0.016, 0.02$.\\label{zeta}} \\end{figure}   The Helium abundance in such a model could be estimated. To a first approximation, \\begin{eqnarray} Y=&&(0.2378+0.0073\\ln\\eta_{10})(1-0.058/\\eta_{10})\\\\ &&+0.013(N_{\\nu}-3)+2\\times 10^{-4} (\\tau_n - 887). \\end{eqnarray} The speed up factor \\begin{equation} \\zeta \\equiv H(a)/H_{\\mathrm{sm}}(a) \\end{equation} is related to the neutrino number by \\begin{equation} \\zeta^{2}=1+\\frac{7}{43}(N_{\\nu}-3). \\end{equation} So we have \\begin{equation} \\Delta Y = 0.08(\\zeta^2 -1). \\end{equation}  The differential speed up factor $\\zeta-1$ for a number of models is plotted in Fig.~\\ref{zeta}. As an example, let us consider $\\alpha=10$, $\\xi=0.004$, and $Q_0 =5.5$ which satisfies the solar system limit $|\\xi| <0.022 Q_{0}^{-1}$. The helium abundance could be reduced by as much as $0.096\\%$. In Fig.~\\ref{Y-eta}, the helium abundance for this NMC model is plotted. It lies much more comfortably within the allowed range. Alternatively, the helium bound on neutrino number could be relaxed. If we apply this to the current cosmological limit on neutrino number \\cite{PDG}, which is $1.7<N_{\\nu}<4.3$ at $95\\%$ C.L., the upper limit of $N_{\\nu}$ could be lifted to 5.  In summary, I have shown that in some NMC models, the BBN helium abundance could be reduced, thus alleviate the marginal disagreement between theory and observation, and make more room for new neutrinos or other new particles. Whether such a reduction is necessary depends on the result of future observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011019_arXiv.txt": {
        "abstract": "Based on detailed models for the explosions, light curves and NLTE-spectra, evolutionary effects of Type Ia Supernovae (SNe~Ia) with redshift have been studied to evaluate their size on cosmological time scales, how the effects can be recognized and how one may be able to correct for them.  In the first part, we summarize the current status of scenarios for Type Ia Supernovae, including the explosion of a Chandrasekhar mass           white dwarf ($M_{Ch}$-WD), the merging scenario and the helium-triggered explosions of low-mass WDs. We show that delayed detonation models  can account for the majority of observations of spectra and light curves. IR observations are a new and powerful tools  to constrain explosion models and details of the flame propagation in the WD. A strong Mg II line at 1.05 $\\mu m$ shows that nuclear burning takes place at the outer, low density layers. This requires a transition from the deflagration to the detonation regime of the nuclear burning front, or a very fast deflagration close to the speed of sound. We put the models into context with the empirical brightness decline relation which is widely applied to use SNe~Ia as yardsticks on cosmological distance scales. This relation can be well understood in the framework of $M_{Ch}$-WDs as a consequence of the opacity effects in combination with the amount of $^{56}Ni$ which determines the brightness. However, the amount of $^{56}Ni$ actually produced depends on a combination of free parameters such as central density and the chemical composition of the WD, and  the propagation of the burning front. We get a spread of $\\approx 0.4 ^m$ around the mean relation which is  larger than currently favored by observations ($\\approx 0.12^m$, \\cite{riess98}) which may hint of an underlying coupling of the progenitor, the accretion rates and the propagation of the burning front.  In a second  part, we investigate the possible evolutionary effects in SNe~Ia  both with respect to changes in the sample of SNe~Ia and individual variations, and how they can be identified. We find that evolution may produce an offset in the brightness decline relation  is restricted to a few tenth of a magnitude. The effects reveal themself by  changes in the U and UV fluxes, and in a change in the maximum brightness/decline relation by $\\Delta M \\approx 0.1 \\times \\Delta t$ where $\\Delta t$ is the difference between local and distant SN-samples. According to new data \\cite{aldering2000}, $\\Delta t \\leq 1  d$ and, thus, evolution is unlikely to eliminate a need for $\\Lambda $. ",
        "introduction": "Two of the important developments in observational SN-research in the last few years were to establish the long-suspected correlation between the peak brightness of SNe~Ia and their rate of decline by means of modern CCD photometry \\cite{phillips87}, and the exact distance calibrations provided by an HST Key Project (e.g. \\cite{saha97}). Both allowed a empirical determinations of $H_o$ with unprecedented accuracy \\cite{hamuy96}. Independent from these calibrations and empirical relations, $H_o$ has been determined by   comparisons of detailed theoretical models for light curves and spectra with observations \\cite{mh94,hk96,nugent96}. All determinations of the Hubble constant are in good agreement with one another. Recently, the routine successful detection of SNe at large redshifts, z  \\cite{perlmutter98,riess98}, has provided an exciting new tool to probe cosmology.  This work has provided results that are consistent with a low matter density in the Universe and, most intriguing of all, yielded hints for a positive cosmological constant. The cosmological results rely on  empirical brightness-decline relations which are calibrated locally. One of the main concerns are systematic, evolutionary effects in the properties of SNe~Ia. Both from  theory and observations, there are some hints for the presents of evolution.  Observational evidence include the finding of Branch et al. \\cite{branch98} who have shown that the mean peak brightness is dimmer in ellipticals than in spiral galaxies.  Wang, H\\\"oflich \\& Wheeler \\cite{wang97}  found that the peak brightness in the outer region of spirals is similar to those found in ellipticals, but that in the central region both intrinsically brighter and dimmer SNe~Ia occur. This implies that the progenitor populations are more inhomogeneous in the inner parts of spirals which contain both young and old progenitors.  From theory,  time evolution is expected to produce the following main effects: (a) a lower   metallicity will decrease the time scale for  stellar evolution of individual stars by about 20 \\% from Pop I to Pop II stars \\cite{schaller92} and, consequently, the  progenitor population which contributes to the SNe~Ia rate at any given time. The stellar radius also shrinks. This will influence the statistics of systems with mass overflow; (b) early on, we expect that systems with shorter life time will dominate the sample whereas, today, old system are contributing which may have not occurred early on. In addition, some scenarios with a life time comparable to the age of the univers such as two merging WDs may be absent a few Gyrs ago, but they may contribute today. (c) The initial metallicity will determine the electron to nucleon fraction of the outer layers  and hence affects the products of nuclear burning; (d) Systems with a shorter life time may dominate early on and, consequently, the typical C/O ratio  of the central region of the WD will be reduced; (e) The properties of the interstellar medium may change; (f) The influence of Z on the structure of WDs may change, but this effect has remains very small (e.g. $2\\% $ when comparing solar with 0.01 solar, \\cite{hwt98}) (g) The distribution of C and O will depend on Z as it influences the `normal' stellar evolution and the properties of the C/O core \\cite{schaller92} (h) The metallicity will effect nuclear burning during the accretion phase of the progenitor \\cite{nomoto2000}.  Based on theoretical models described in \\S III, we want to study the possible effects of the change of the progenitor and its metallicity on the light curves and spectra of SNe~Ia. Note that most of the results have been obtained in serveral collaborations over the years (see acknowledgments).  \\noindent ",
        "conclusions": ""
    },
    "0011/hep-th0011110_arXiv.txt": {
        "abstract": "These notes provide a brief introduction to modern cosmology, focusing primarily on theoretical issues.  Some attention is paid to aspects of potential interest to students of string theory, on both sides of the two-way street of cosmological constraints on string theory and stringy contributions to cosmology.  Slightly updated version of lectures at the 1999 Theoretical Advanced Study Institute at the University of Colorado, Boulder. ",
        "introduction": "\\label{sec:introduction}  String theory and cosmology are two of the most ambitious intellectual projects ever undertaken.  The former seeks to describe all of the elements of nature and their interactions in a single coherent framework, while the latter seeks to describe the origin, evolution, and structure of the universe as a whole.  It goes without saying that the ultimate success of each of these two programs will necessarily involve an harmonious integration of the insights and requirements of the other.  At this point, however, the connections between cosmology and string theory are still rather tenuous.  Indeed, one searches in vain for any appearance of ``cosmology'' in the index of a fairly comprehensive introductory textbook on string theory \\cite{polchinski}, and likewise for ``string theory'' in the index of a fairly comprehensive introductory textbook on cosmology \\cite{peebles}.  These absences cannot be attributed to a lack of knowledge or imagination on the part of the authors.  Rather, they are a reflection of a desire to stick largely to those aspects of these subjects about which we can speak with some degree of confidence (although in cosmology, at least, not everyone is so timid \\cite{kt,lindebook}). In cosmology we have a very successful framework for discussing the evolution of the universe back to relatively early times and high temperatures, which however does not reach all the way to the Planck era where stringy effects are expected to become important.  In string theory, meanwhile, we have learned a great deal about the behavior of the theory in certain very special backgrounds, which however do not include (in any obvious way) the conditions believed to obtain in the early universe.  Fortunately, there is reason to believe that this situation may change in the foreseeable future.  In cosmology, new data coming in from a variety of sources hold the promise of shedding new light on the inflationary era that is widely believed to have occurred in the early universe, and which may have served as a bridge from a quantum-gravity regime to a classical spacetime. And in string theory, the last few years have witnessed a number of new proposals for formulating the theory in settings which were previously out of reach, and there are great hopes for continued progress in this direction. Furthermore, there is a reasonable expectation of significant improvement in our understanding of particle physics beyond the standard model, from upcoming accelerator experiments as well as attempts to directly detect cosmological dark matter.  It is therefore appropriate for cosmologists and string theorists to keep a close watch on each other's work over the next few years, and this philosophy has guided the preparation of these lectures.  I have attempted to explain the basic framework of the standard cosmological model in a mostly conventional way, but with an eye to those aspects which would be most relevant to the application of string theory to cosmology.  (Since these lectures were delivered, several reviews have appeared which discuss aspects of string theory most relevant to cosmology \\cite{lwc,Banks:1999ay,Veneziano:2000pz,Easson:2000mj}.) My goals are purely pedagogical, which means for example that I have made no real attempt to provide an accurate historical account or a comprehensive list of references, instead focusing on a selection of articles from which a deeper survey of the literature can be begun.  Alternative perspectives can be found in a number of other recent reviews of cosmology \\cite{Lyth:1993eu,Kolb:1994ad, Brandenberger:1996kf,Ellis:1998ct,turnertyson,Lazarides:1998pz}.  (Note:  These lectures were first written and delivered in summer 1999.  I have added occasional references to subsequent developments where they seemed indispensible, but have made no effort at a thorough updating.) ",
        "conclusions": "\\label{sec:conclusions}  The last several years have been a very exciting time in string theory, as we have learned a great deal about non-perturbative aspects of the theory, most impressively the dualities connecting what were thought to be different theories. They have been equally exciting in cosmology, as a wealth of new data have greatly increased our knowledge about the constituents and evolution of the universe.  The two subjects still have a long way to go, however, before their respective domains of established understanding are definitively overlapping.  One road toward that goal is to work diligently at those aspects of string theory and cosmology which are best understood, hoping to enlarge these regions until they someday meet.  Another strategy is to leap fearlessly into the murky regions in between, hoping that our current fumbling attempts will mature into more solid ideas.  Both approaches are, of course, useful and indeed necessary; hopefully these notes will help to empower the next generation of fearless leapers."
    },
    "0011/astro-ph0011527_arXiv.txt": {
        "abstract": "We describe the discovery circumstances and photometric properties of 2000 EB173, now one of the brightest trans-Neptunian objects (TNOs) with opposition magnitude $m_{R}=18.9$ and also one of the largest Plutinos, found with the drift-scanning camera of the QUEST Collaboration, attached to the 1-m Schmidt telescope of the National Observatory of Venezuela. We measure  $B-V=0.99 \\pm 0.14$ and $V-R=0.57 \\pm 0.05$, a red color observed for many fainter TNOs. At our magnitude limit $m_{R}=20.1 \\pm 0.20$, our single detection reveals a sky density of 0.015 (+0.034, -0.012) TNOs per deg$^{2}$ (the error bars are $68\\%$ confidence limits), consistent with fainter surveys showing a cumulative number proportional to $10^{ 0.5 m_{R} }$.  Assuming an inclination distribution of TNOs with FWHM exceeding 30 deg, it is likely that one hundred to several hundred objects brighter than $m_{R}=20.1$ remain to be discovered. ",
        "introduction": "In recent years, the existence of a belt of small planets with orbits beyond  Neptune has been firmly established following the discovery of 1992 QB1 (Jewitt \\& Luu 1993). These  bodies are believed to be icy planetesimals remaining from the time preceding planet formation, as predicted by Edgeworth (1949) and Kuiper (1951). As of mid-2000, over 300 trans-Neptunian objects (TNOs) are known (Marsden 2000), most having faint magnitudes $(20< m_{R} <28)$ owing to their detection in narrow-field surveys with large-aperture telescopes. Physical studies have been limited to broadband filter photometry for the fainter bodies and spectroscopy for a few of the brighter objects. These studies show that TNOs are spectrally diverse, with colors ranging from neutral to extremely red (Luu \\& Jewitt 1996; Tegler \\& Romanishin 1998; Jewitt \\& Luu 1998; Noll, Luu, \\& Gilmore 2000; Barucci  et al. 2000) and with absorption features indicative of hydrocarbons and water ice (Brown et al. 1997; Brown, Cruikshank, \\& Pendleton 1999). Previous analyses of the discovery data have shown that the number brighter than $m_{R}$ increases as $10^{ \\alpha m_{R}}$, with $ \\alpha $ ranging from $0.52 \\pm 0.02$ to $0.76 \\pm 0.10$ (Jewitt, Luu \\& Trujillo 1998; Gladman et al. 1998; Chiang \\& Brown 1999) depending on the choice and interpretation of discovery data. It is not known if this luminosity function extends to magnitudes as bright as Pluto ($m_{V}=14$) or instead cuts off at intermediate magnitudes. Such a cutoff might reveal a termination in the growth phase of the TNOs at the time of Neptune's formation (Bailey 1993; Stern \\& Colwell 1997).  Here we describe the discovery circumstances and photometric properties of 2000 EB173, now one of the brightest TNOs with opposition magnitude $ m_{R}=18.9$ and also one of the largest Plutinos (a TNO with a stable orbit like Pluto completing two orbits for every three by Neptune, thereby avoiding close encounters with Neptune). We show that this detection would not be likely unless the luminosity function continues with a shallow slope ($ \\alpha \\sim 0.5$) at bright magnitudes, $m_{R} < 19$.  We also show that the object's photometric properties are similar to those of fainter TNOs with red color. ",
        "conclusions": "1. The unusual brightness of 2000 EB173 provides an opportunity to study the physical nature of one TNO with substantially greater precision than has been possible previously. It now becomes possible for all astronomers, including those with modest facilities, to make a valuable contribution to the study of the primordial bodies preserved in the outer solar system.  Possible contributions would be to determine the size of 2000 EB173 by measuring the albedo with infrared observations or by direct observation with space-based telescopes. Such studies would have important implications for the total mass of the TNO population, currently uncertain by more than an order of magnitude.  2. We observe no variability in the lightcurve of 2000 EB173 and no cometary coma. The color is very red shortwards of 0.7 microns, typical of many fainter TNOs. Our measurements indicate a possible flattening of the reflectance longwards of 0.7 microns, but are also consistent with a redder color.  3. Our detection of one TNO brighter than $m_{R}=20.1$ in a search of 66.8 deg$^{2}$ favors the flat luminosity function ($ \\alpha = 0.58$) of Jewitt, Luu, \\& Trujillo (1998) over the steeper distribution ($ \\alpha = 0.768$) of Gladman et al. (1998).  4. Extrapolating to a search of the entire area within 60 deg of the ecliptic (within the inclination range for TNOs determined by Jewitt, Luu, \\& Chen 1996), one could reasonably expect to find from one hundred to several hundred TNOs brighter than $m_{R}=20.1$. A dedicated search using a telescope with  capabilities similar to our own would discover most of these bodies in a few years."
    },
    "0011/astro-ph0011461_arXiv.txt": {
        "abstract": "X-ray galaxy clusters are ideal tracers of the large-scale structure of the Universe. Based on the ROSAT All-Sky Survey we are constructing a sample of the X-ray brightest galaxy clusters for cosmological studies. We have already completed the  compilation of a sample of about 1000 clusters in the sky excluding a 40 degree wide band around the galactic plane. We have demonstrated  with the southern sample, which is the so far most complete, that we can get reliable statistical measures of the large-scale structure of the Universe to scales in excess of $400 h^{-1}$ Mpc. With the experience gained in these projects we are extending the search for X-ray galaxy clusters into the Zone-of-Avoidance. We have compiled a sample of 181 good cluster candidates with a well defined completeness of about 70\\% down to a flux limit of $3\\cdot 10^{-12}$ erg s$^{-1}$ in the zone with an absorption depth of $A_v \\le 2.25$. The candidates are definitely identified and redshifts are obtained in an ongoing optical follow-up program. ",
        "introduction": "\\begin{figure}[h] \\plotone{Boehringer1.eps} \\caption{Sky distribution of the X-ray cluster surveys based on the RASS in galactic coordinates. The solid data points show the current samples of the NORAS Survey (northern sky) and REFLEX Survey (southern sky up to declination $d \\le 2.5\\deg$). Also shown are the cluster candidates found in the Zone-of Avoidance $|b_{II}| \\le 20\\deg $ and $A_v \\le 2.5$ (open circles).} \\end{figure}    Clusters of galaxies as the largest well defined building blocks of our Universe are ideal probes for the study of the cosmic large-scale structure and for the test of cosmological models (e.g. Henry et al. 1992, Dalton et al. 1994, Collins et al. 1995, Eke, Cole \\& Frenk 1996, Borgani et al. 1999). Statistical measures of the galaxy cluster population like the cluster mass function, the two-point-correlation function, and the density fluctuation power spectrum can give very important constraints on the characteristic measures of the matter density distribution throughout the Universe and its evolution as a function of time. For such studies the compilation of a statistically complete cluster catalogue with a well controlled selection function is the important first step.  Ideally one would like to compile a cluster catalogue where the cluster are selected by their gravitational mass. The mass is not a direct observable, however, and one relies on relations of observable quantities that are closely correlated to the cluster mass. While cluster catalogues have first been compiled on the basis of optical photographic survey data (e.g. Abell 1958, Abell et al. 1989) it has been recognized that it is very difficult to obtain a cluster mass from the observed optical ``richness'' parameter. X-ray observations provide a much better means for the construction of cluster catalogues for the study of the large-scale structure. The X-ray emission observed from clusters originates from the thermal emission of hot intracluster gas bound in the gravitational potential well of the clusters (e.g. Sarazin 1986). Therefore the plasma emission is a certain indication of a true gravitationally bound structure. In addition the thermal emission for the typical intracluster plasma temperatures of several keV has the radiative emission maximum in soft X-rays were the available X-ray telescopes are most effective. This makes galaxy clusters readily detectible out to large distances with present X-ray telescopes. The main advantage of the X-ray detection is, however, the fact that the X-ray luminosity is closely correlated to the cluster mass (Reiprich \\& B\\\"ohringer 1999) with a dispersion of about 50\\% in the determination of the mass for a given X-ray luminosity. ",
        "conclusions": "To illustrate the efficiency of the cluster detection as a function of the interstellar absorbing column density, $n_H$, we show in Fig. 7 the surface density of the cluster counts in the sky as a function of $n_H$. Also shown is the surface density of all REFLEX survey clusters and those which are found to be extended. In the low $n_H$ region up to about $n_H \\le 2.0 \\cdot 10^{21}$ cm$^{-2}$ we almost recover the surface density of the extended REFLEX clusters which is almost 80\\% of the total sample. Beyond $n_H \\sim 2.0 \\cdot 10^{21}$ cm$^{-2}$ the detection efficiency drops significantly showing that the cluster search becomes more difficult. In summary we estimate that we can recover at least 60 - 70\\% of the clusters above a flux limit of $3 \\cdot 10^{-12}$ s$^{-1}$ cm$^{-2}$ in 2/3 of the ZoA region which is less affected by absorption. Definite identifications and redshifts for these cluster candidates are currently obtained in an ongoing follow-up observation program in the northern and southern sky."
    },
    "0011/hep-ph0011318_arXiv.txt": {
        "abstract": "s{Two germanium detectors are currently operating in the Canfranc Underground Laboratory at 2450 m.w.e looking for WIMP dark matter. One is a 2 kg $^{76}$Ge IGEX detectors (RG-2) which has an energy threshold of 4 keV and a low-energy background rate of about 0.3 c/keV/kg/day. The other is a small (234 g) natural abundance Ge detector (COSME), of low energy threshold (2.5 keV) and an energy resolution of 0.4 keV at 10 keV which is looking for WIMPs and for solar axions. The analysis of 73 kg-days of data taken by COSME in a search for solar axions via their photon Primakoff conversion and Bragg scattering in the Ge crystal yields a 95\\% C.L. limit for the axion-photon coupling $g_{a \\gamma \\gamma} < 2.8 \\times 10^{-9}$  GeV$^{-1}$. These data, analyzed for WIMP searches provide an exclusion plot for WIMP-nucleon spin-independent interaction which improves previous plots in the low mass region. On the other hand, the $\\sigma(m)$ exclusion plot derived from the 60 kg-days of data from the RG-2 IGEX detector improves the exclusion limits derived from other ionization (non thermal) germanium detector experiments in the region of WIMP masses from 30 to 100 GeV recently singled out by the reported DAMA annual modulation effect.} ",
        "introduction": "Substantial evidence exists suggesting most matter in the universe is dark, and there are compelling reasons to believe it consists mainly of non-baryonic particles. Among these candidates, Weakly Interacting Massive Particles (WIMPs) and Axions are among the front runners. The lightest stable particles of supersymmetric theories, like the neutralino, describe a particular class of WIMPs\\cite{Gri}.  Direct detection of WIMPs rely on the measurement of their elastic scattering off target nuclei in a suitable detector\\cite{Mor99}. Slow moving ($\\sim300$ km/s) and heavy ($10 - 10^3$ GeV) galactic halo WIMPs could make a Ge nucleus recoil with a few keV, at a rate which depends on the type of WIMP and interaction. Ultra--low background detectors with very low energy thresholds are needed for such purpose. Germanium detectors have reached one of the lowest background levels of any type of detector and  have a reasonable ionization yield($\\sim 0.25$). Thus, with sufficiently low energy thresholds, they are attractive devices for WIMP direct detection. In addition, solar axions can also been searched \\cite{Paschos:1994yf,Creswick:1998pg} with these crystal detectors by looking for the Primakov conversion into photons and Bragg diffraction.  Two experiments with germanium detectors are currently running at the Canfranc Underground Laboratory: IGEX-DM and COSME-2. Data obtained with both experiments are presented in this work, as well as the WIMP exclusions obtained from the analysis of both sets of data. COSME data has also been analyzed looking for the solar axion signal mentioned above. ",
        "conclusions": ""
    },
    "0011/astro-ph0011131_arXiv.txt": {
        "abstract": "We present models of interacting galaxies in order to study the development of spatial offsets between the gaseous and stellar components in tidal tails. Observationally, such offsets are observed to exist over large scales (\\eg NGC 3690; Hibbard \\etal 2000), suggesting an interaction between the tidal gas and some (unseen) hot ISM. Instead, our models show these offsets are a natural consequence of the radially extended \\HI spatial distribution in disk galaxies, coupled with internal dissipation in the gaseous component driven by the interaction. This mechanism is most effective in systems involved in very prograde interactions, and explains the observed gas/star offsets in interacting galaxies without invoking interactions with a hot ISM, starburst ionization, or dust obscuration within the tails. ",
        "introduction": "Since the original work of Toomre \\& Toomre (1972) and Wright (1972), it has been recognized that the development of tidal tails during galaxy encounters is predominantly a kinematic effect -- stars are tidally torn away from the parent galaxy and orbit as ``test particles'' in the evolving potential. While self-gravity and hydrodynamic forces can play a significant role in the development of substructure {\\it within} tidal debris (see, \\eg Wallin 1990; Barnes \\& Hernquist 1992; Elmegreen, Kaufman, \\& Thomasson 1993), the {\\it global} evolution of tidal tails is much simpler.  Because both gas and stars respond to the gravitational potential in a similar fashion, they should have similar kinematics and spatial morphology on large scales.  To be sure, differences exist between the mass distribution of stars and gas in tidal tails -- gaseous tails are often much more extended than their stellar counterparts -- but this is due largely to the differing radial distribution of stars and gas in the parent galaxy. Because \\HI disks extend further out than stellar disks (\\eg Broeils \\& van Woerden 1994; Broeils \\& Rhee 1997), there is more loosely bound gas to be ejected to large distances in the tidal debris. Nonetheless, the kinematic nature of the process predicts that when material is tidally ejected, stars and gas should remain largely cospatial in the resulting tidal tails.  Observationally, this situation holds true in many cases. Examples of cospatial gas/star tails include nearly all the nearby mergers such as the Antennae, NGC 7252, or NGC 4676 (Hibbard \\& van Gorkom 1996). This coevolutionary status of tidal gas and stars permits detailed matching between the theoretic predictions of stellar-dynamical models and the observed morphology and kinematics of the gas. Several investigators have taken advantage of this fact to produce dynamical models of specific merging systems (\\eg Stanford \\& Balcells 1991; Mihos \\etal 1993; Hibbard \\& Mihos 1995; Yun 1999; Barnes \\& Hibbard 2000).  However, in a significant number of systems the \\HI and stellar tidal debris are not purely cospatial.  In many of these cases, the observed gas/star offsets can be explained, at least qualitatively, by interpenetrating events during galaxy collisions. For example, the offset in the tidal bridge of NGC 7714/5 (Smith \\etal 1997) is reasonably expected to occur since tidal bridges form from material on the interface between two interacting galaxies; collisional shocks will act to strip gas material out of the developing tidal bridge. In retrograde encounters, short tails and plumes form from material which passes entirely through the companion galaxy, leading to a strong decoupling of the stars and gas (compare, for example, Figures 11 and 12 of Mihos \\& Hernquist 1996). In yet other cases, offsets can be attributed to the fact that tidal material has fallen back through the the merging galaxies and may be subject to additional dissipation, as perhaps seen in the offset stellar and gaseous ``shells'' in Centaurus A (Schiminovich \\etal 1994).  However, these explanations cannot account for a small number of cases in which offsets exist between the gas and stars in extended tidal tails. Notable examples include the tidal tails of NGC 520 (Hibbard, Vacca, \\& Yun 2000), 7714/5 (Smith \\etal 1997), and, most dramatically, NGC 3690 (Hibbard \\& Yun 1999), in which the gas and stars show an offset of $\\sim$ 20 kpc over much of the $\\sim$ 180 kpc extent of the tail.  In these examples, none of the previously mentioned mechanisms easily explains the observed offsets. These are all examples of long tidal tails, which develop from the side of the galaxy {\\it away} from the companion and do not therefore come in physical contact with the companion. Unlike the tidal shells and loops in older mergers, the long tails here have not yet fallen back through their central objects, and, furthermore, the length of the tails precludes a retrograde geometry for the encounter.  At face value, these observations call into question the idea that the evolution of tidal tails is purely kinematic, and suggest that hydrodynamic forces may also play a role in the global evolution of the tidal tails. Hibbard \\etal (2000) suggest that the offsets may arise from a variety of effects, including ionization of part of the tail from a starburst wind or from the stars in the tail itself, partial obscuration by dust in the tail, or by an interaction between the tidal gas and a hot ISM component, either a hot gaseous halo or a starburst wind. However, many of these solutions seem somewhat contrived, particularly with respect to the case of NGC 3690. In this case, the gas appears to lead the stellar tail, whereas a gas tail expanding into a hot halo would tend to trail its stellar counterpart. A starburst wind solution is also difficult to envision, as the wind would need to be extremely collimated over very large scales to result in such an extended, linear offset. While ROSAT observations of NGC 3690 (Heckman \\etal 1999) do show evidence for a hot outflow, but the X-ray isophotes exist on scales much smaller than needed to explain the offset tidal gas ($\\sim$ 45 kpc compared to the 180 kpc extent of the tidal tail).  While the starburst wind may extend to even larger radii at lower temperatures and densities, there is no direct evidence for such an extended outflow.  In this paper we describe a simpler solution for the observed gas/star offsets in some tidal tails. In this model, the offsets arise from the differing radial distributions of gas and stars, coupled with interaction-driven dissipation and inflow in the tidal gas. This scenario explains the observed offsets without resorting to esoteric processes such as hot ISM interactions, starburst ionization, or dust obscuration. ",
        "conclusions": "The fact that we can account for the offset between the gas and stars in tidal tails without the need for a separate hot ISM component argues against these offsets necessarily arising from interactions with gaseous halos or starburst winds. We emphasize here that the initial conditions are far from being contrived; on the contrary, the radial distribution of \\HI in galaxies is typically much more extended than the stellar component (\\eg Broeils \\& van Woerden 1994; Broeils \\& Rhee 1997).  In truth, the extended gas disk in our model is a more realistic initial galaxy model than those used by MH96 or Barnes \\& Hernquist 96, where purely exponential gas distributions were used. In essence, better initial conditions have yielded a more accurate description of the formation and evolution of the tidal tails, showing the natural development of the gas/star offset.  These models indicate that the offset is produced when the tail comes from a very prograde encounter; if the encounter is fairly inclined, the dissipation in the gas component is reduced and the offset inhibited. How does this prediction relate to the observed \\HI properties of mergers? Most nearby mergers do not show strong offsets, consistent with the idea that the geometry must be fairly specific to yield offsets. In two of the more extreme examples of tidal offsets -- NGC 3690 and NGC 520 -- the orbital geometry inferred from the measured kinematics of the system suggests a prograde encounter (Hibbard \\etal 2000), although with significant uncertainties in the exact geometry. If our explanation of the offset holds true, our constraints on the geometry are even more severe than those inferred from the global kinematics of these systems.  Our model also explains the fact that the offset gas tail in NGC 3690 appears to curve back onto the optical tail at large radius (the ``hook'' described by Hibbard \\& Yun 1999). Material in the tail at large radius comes from the most loosely bound material at initially large radius in the progenitor disk. Coming from the outskirts of the disk,  this material is less susceptible to the inflow driven by the self-gravitating response of the host stellar disk. This material develops in an almost purely collisionless evolution so that the stars and gas remain cospatial. It is only further inwards along the tail that the offset develops, since it is here where the tidal material would have come from deeper in the progenitor galaxy, and be more prone to dissipative effects.  Given the different spatial distributions of gas and stars in the tidal tails, it is also interesting to ask if we see kinematic segregation between the components as well. This has particularly important ramifications for the frequent exercise of modeling specific galaxies based on \\HI or H$\\alpha$ velocity fields. Because of the increased computational cost of running hydrodynamic simulations, collisionless N-body models are typically used to match the observed (gas-phase) velocities (\\eg Stanford \\& Balcells 1991; Mihos \\etal 1993; Hibbard \\& Mihos 1995; Yun 1999; Barnes \\& Hibbard 2000).  If severe kinematic segregation exists in the tidal debris, models derived from matching collisionless kinematics to gas velocities may be in serious error.  To examine this possibility, Figure 6 shows a simulated position-velocity plot for both the gas and stars, observed in the plane of the prograde tail such that the portion of prograde tail which shows the offset moves directly towards the observer.  This viewing geometry results in the {\\it maximum} observed kinematic offset between the stars and gas in the prograde tail.  These position-velocity plots show that the kinematic offsets are observable, with offset amplitudes of $\\sim$ 50--70\\kms, but that they do not radically affect the inferred global kinematics of the system. While detailed model-data matching may hint at differences between the stellar and gas kinematics (see, for example, the discussion in Hibbard \\& Mihos 1995), the derived models should not be seriously misled by these effects.  Finally, we emphasize that our explanation of gas/star offsets in tidal debris is not the only possible mechanism which may produce offsets. Certainly many examples exist of kinematic decoupling which occurs when the encounter draws tidal material through the inner portions of galaxies, such as retrograde encounters or late infall from tidal tails. The anti-correlation noted by Hibbard \\etal (2000) between gas and stars in the tidal debris around Arp 220 are one possible example of this situation, where the stubby tidal features seen in the optical are most likely not due to an extremely prograde encounter. Also, as Hibbard \\etal (2000) point out in their discussion of gas/star offsets, many interacting galaxies possess strong starburst winds (Heckman, Armus, \\& Miley 1987, 1990) and models suggest that collisionally heated halo gas may exist (Barnes \\& Hernquist 1996); the interaction between the tidal \\HI and these hot gas components may yet be significant in some cases. What our model shows is that the observed offsets in long tidal tails are easily explained without the need for a hot component, but they do not rule out the presence -- or even the effects -- of such a hot component."
    },
    "0011/astro-ph0011307_arXiv.txt": {
        "abstract": "{ The observation of the supernova remnant N132D by the scientific instruments on board the \\xmm\\ satellite is presented. The X-rays from N132D are dispersed into a detailed line-rich spectrum using the Reflection Grating Spectrometers. Spectral lines of C, N, O, Ne, Mg, Si, S, and Fe are identified. Images of the remnant, in narrow wavelength bands, produced by the European Photon Imaging Cameras reveal a complex spatial structure of the ionic distribution. While K-shell Fe emission seems to originate near the centre, all of the other ions are observed along the shell. A high O$^{6+}$ / O$^{7+}$ emission ratio is detected on the northeastern edge of the remnant. This can be a sign of hot ionising conditions, or it can reflect relatively cool gas. Spectral fitting of the CCD spectrum suggests high temperatures in this region, but a detailed analysis of the atomic processes involved in producing the O$^{6+}$ spectral lines leads to the conclusion that the intensities of these lines alone cannot provide a conclusive distinction between the two scenarios. ",
        "introduction": "N132D, one of the brightest soft X-ray sources in the Large Magellanic Cloud (LMC), was identified as a Type II supernova remnant (SNR) by Westerlund \\& Mathewson\\ (\\cite{westerlund66}). The high velocities of the oxygen-rich (O-rich) ejecta near the centre of the remnant were first detected by Danziger \\& Dennefeld\\ (\\cite{danziger76}) and later measured precisely by Morse et al.\\ (\\cite{morse95}). The fast moving ejecta is believed to originate from a core collapse supernova (SN) of a massive star. Blair et al.\\ (\\cite{blair00}) have recently performed a thorough optical and UV study of N132D using the Hubble Space Telescope (HST). With the exceptionally high spatial resolution of the HST, they were able to distinguish between SN ejecta and swept-up interstellar matter (ISM). Observing a knot on the northwestern rim of the remnant, they found elemental abundances that agree fairly well with those previously obtained from X-ray observations of the entire remnant (Hughes et al.\\ \\cite{hughes98}; Favata et al.\\ \\cite{favata98}), and which are also consistent with the mean LMC abundances. In the ejecta, on the other hand, only C, O, Ne, and Mg could be detected in the optical/UV spectra. The absence of O-burning elements, such as Si, S, and Fe, in the ejecta has led Blair et al.\\ (\\cite{blair00}) to suggest that N132D may be a product of a Type Ib SN explosion.  N132D was first observed in the X-rays with the \\einstein\\ Observatory (Mathewson et al.\\ \\cite{mathewson83}). Using the spectrometers on board \\einstein, Hwang et al.\\ (\\cite{hwang93}) detected individual emission lines of O$^{7+}$, O$^{6+}$, Ne$^{9+}$, Ne$^{8+}$, and Fe$^{16+}$. Comparing with the \\rosat\\ images, Morse et al.\\ (\\cite{morse95}) found little overlap between the optical and soft X-ray emitting regions. The X-ray emission seemed to come from a large shell of shocked ISM gas (see also Williams et al.\\ \\cite{williams99}), which is also the source of strong synchrotron radio emission (Dickel \\& Milne\\ \\cite{dickel95}). Hughes et al.\\ (\\cite{hughes98}) fitted the \\asca\\ CCD spectrum of N132D, obtaining mean LMC abundances, which further suggested that the X-rays in the remnant are predominantly emitted by shocked ISM.  In this work, we present the X-ray spectrum of N132D measured with the Reflection Grating Spectrometers (RGS), complemented by the images taken with the European Photon Imaging Cameras (EPIC), both on board \\xmm. Owing to the uniquely high dispersion of the RGS, almost all of the spectral features emitted by this extended ($\\sim$120\\arcsec) source are resolved. Along with the RGS observation of 1E0102.2$-$7219 in the Small Magellanic Cloud (Rasmussen et al.\\ \\cite{andy01}), these are the first highly resolved X-ray spectra of extended SNRs. ",
        "conclusions": "The RGS spectrum of N132D features lines of C, N, O, Ne, Mg, Si, S, and Fe. K-shell Ar, Ca, and Fe are detected in the \\pn\\ spectrum. The narrow-band EPIC images of N132D give a coarse mapping of the elemental and temperature structure in the remnant. With the exception of O, the dominant part of the soft X-ray emission originates from shocked ISM along the southeastern, and northwestern edges of the expanding shell. In contrast, strong Fe-K emission is detected near the centre of the remnant, perhaps indicating that high-Z ejecta is too highly ionised to be observed in shorter wavelengths. O$^{7+}$ is present throughout the remnant, while the brightest spot of O$^{6+}$ emission is near the northeastern rim. The O$^{6+}$ emission can be attributed, either to low temperatures, or to hot, recently shocked material that is in the process of being ionised. Fitting of the \\pn\\ spectrum indicates high temperatures in that region. Unfortunately, the emission lines of O$^{6+}$ alone, although clearly seen in the RGS spectrum, cannot provide a conclusive distinction between these two scenarios, due to considerations associated with the atomic processes involved in producing these lines. A more detailed analysis of the RGS spectrum, and especially of the many strong Fe lines, will potentially provide a better insight into the temperature and non-equilibrium structure of the hot gas in N132D."
    },
    "0011/astro-ph0011077_arXiv.txt": {
        "abstract": "s{ There is growing evidence that, at the $S_{850\\mu m} < 1$mJy level, the sub-mm galaxy population (and hence a potentially significant fraction of the sub-mm background) is associated with the star-forming Lyman-break population already detected at optical wavelengths. However, the implied star-formation rates in such objects (typically $3 - 30 {\\rm M_{\\odot} yr^{-1}}$) fall one or two orders of magnitude short of the level of star-forming activity required to produce the most massive elliptical galaxies on a timescale $\\sim 1$Gyr. If a significant fraction of massive ellipticals did form the bulk of their stars in short-lived massive starbursts at high redshift, then they should presumably be found among the brighter, $S_{850\\mu m} \\simeq 10$mJy sub-mm sources which are undoubtedly {\\it not} part of the Lyman-break population. A first powerful clue that this is indeed the case comes from our major SCUBA survey of radio galaxies, which indicates that massive dust-enshrouded star-formation in at least this subset of massive ellipticals is largely confined to $z > 2.5$, with a mean redshift $z \\simeq 3.5$. While radio selection always raises concerns about bias, I argue that our current knowledge of the brightest ($S_{850\\mu m} \\simeq 10$mJy) sub-mm sources detected in unbiased SCUBA imaging surveys indicates that they are also largely confined to this same high-$z$ regime. Consequently, while the most recent number counts imply such extreme sources can contribute only 5-10\\% of the sub-mm background, their comoving number density (in the redshift band $3 < z < 5$) is $\\simeq 1 - 2 \\times 10^{-5} {\\rm Mpc^{-3}}$, sufficient to account for the formation of {\\it all} ellipticals of comparable mass to radio galaxies ($\\ge 4L^{\\star}$) in the present-day universe.} ",
        "introduction": "The formation mechanism of elliptical galaxies remains a fundamental and controversial issue in cosmology. In current models of galaxy formation dominated by cold dark matter (CDM), elliptical galaxies arise from the merging at low redshift of intermediate-mass discs \\cite{ba}, and some recent data have been interpreted as supportive of the implied gradual formation of massive ellipticals at relatively low redshift \\cite{kaa}. However, the validity of these analyses has recently been questioned \\cite{to,ji}, and other observational evidence at low/moderate redshift \\cite{bo,no,be}, continues to favour a picture in which at least some massive ellipticals formed the bulk of their stars in a relatively short-lived ($\\le 1$ Gyr) massive starburst at high redshift ($z > 3$) \\cite{ji}. While this high-redshift star-formation scenario might only apply to a subset of ellipticals \\cite{me}, it can be argued that it applies to massive ellipticals in general \\cite{re}. Moreover, the clarification of the link between black-hole and spheroid mass \\cite{ma} suggests that the hosts of AGN may be more representative of spheroids in general than previously supposed.  \\begin{figure} \\vspace{5cm} \\centering \\setlength{\\unitlength}{1mm} \\special{psfile=umassfig1.eps angle=0 hoffset=70 voffset=-30 hscale=32 vscale=32} \\caption{A plot of observed 850-$\\mu m$ flux density against UV SFR for known optically-selected starburst galaxies in the HDF. The bottom panel shows the raw values, while the top panel shows the mean values with standard errors. A statistically significant sub-mm detection is obtained for star-forming galaxies with $SFR > 1 h^{-2} {\\rm M_{\\odot} yr^{-1}}$} \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011288_arXiv.txt": {
        "abstract": "We measured infrared surface brightness fluctuation (SBF) distances to an isotropically-distributed sample of 16 distant galaxies with redshifts reaching 10,000 \\kms\\ using the near-IR camera and multi-object spectrometer (NICMOS) on the {\\it Hubble Space Telescope} (HST). The excellent spatial resolution, very low background, and brightness of the IR fluctuations yielded the most distant SBF measurements to date. Twelve nearby galaxies were also observed and used to calibrate the F160W (1.6 \\micron) SBF distance scale. Of these, three have Cepheid variable star distances measured with HST and eleven have optical $I$-band SBF distance measurements. A distance modulus of 18.5 mag to the Large Magellanic Cloud was adopted for this calibration. We present the F160W SBF Hubble diagram and find a Hubble constant \\Ho$\\,{=}\\,76\\pm 1.3$ (1-$\\sigma$ statistical) $\\pm 6$ (systematic) \\kmsmpc. This result is insensitive to the velocity model used to correct for local bulk motions. Restricting the fit to the six most distant galaxies yields the smallest value of \\Ho$\\,{=}\\,72\\pm 2.3$ \\kmsmpc\\ consistent with the data.  This 6\\% decrease in the Hubble constant is consistent with the hypothesis that the Local Group inhabits an under-dense region of the universe, but is also consistent with the best-fit value of \\Ho$\\,{=}\\,76$ \\kmsmpc\\ at the 1.5-sigma level. ",
        "introduction": "The Hubble constant, \\Ho, is the most fundamental of the cosmological parameters.  Yet in spite of its key role in our understanding of the universe, an accurate determination of its value eluded researchers for decades.  It has only been within the last few years that the promise of knowing \\Ho\\ to better than 10\\% has been realized (Mould et al. 2000). The {\\it Hubble Space Telescope} (HST) has occupied a key role in resolving the debate over the Hubble constant by enabling distance measurements not previously possible from the ground.  With HST's spatial resolution, Cepheid variable stars have been detected in galaxies as distant as 20 Mpc. Cepheid distances to a variety of galaxies, including some in the important Virgo and Fornax clusters, have provided new calibrations of many secondary distance indicators, including type-Ia supernovae (Gibson et al. 2000, Parodi et al. 2000), fundamental plane (Kelson et al. 2000), Tully-Fisher (Sakai et al. 2000), planetary nebulae, globular clusters, and surface brightness fluctuations (Ferrarese et al. 2000a). Uniform HST Cepheid distances were collected by Ferrarese et al. (2000b).  Surface brightness fluctuations have emerged as an accurate and reliable distance indicator (Tonry et al. 1997, Blakeslee et al. 1999). HST has made it possible to not only better calibrate SBFs by providing Cepheid distances to calibration galaxies, but also allowed detection of fluctuations in half a dozen galaxies at much greater distances than possible from the ground (Lauer et al. 1998, Pahre et al. 1999, Thomson et al. 1997). Two additional low signal-to-noise ratio (S/N) measurements in the Coma cluster (Thomson et al. 1997, Jensen et al. 1999), were the most distant SBF measurements until the current NICMOS project.  Surface brightness fluctuations have a much larger amplitude in the near-IR than at optical wavelengths. The Near Infrared Camera and Multi-object Spectrograph (NICMOS) on the HST provides the combination of low background and high spatial resolution needed to measure IR SBFs beyond 100 Mpc for the first time.  The purpose of this study was to calibrate the F160W SBF distance scale and to measure \\Ho\\ beyond the effects of local flows. Reaching distances twice as large as previous SBF studies for a sample uniformly distributed on the sky provided immunity to many of the difficulties that plague all attempts to measure \\Ho\\ within 50 Mpc.  In the next section we describe the selection of the calibration and distant galaxy samples observed. In Section~\\ref{reduxsection} we discuss the procedures used to acquire and reduce the data.  Section~\\ref{sbfsection} describes the methods used to determine the SBF amplitude.  Section~\\ref{calsection} discusses the empirical calibration of the F160W SBF distance scale and the comparison to stellar population models. Section~\\ref{hubblesection} presents the IR SBF Hubble diagram. Finally, we discuss the relationship of our measurement to others which find lower values of \\Ho\\ and conclude with a summary section. ",
        "conclusions": "We measured accurate IR SBF distances to a collection of 16 uniformly-distributed distant galaxies for the purpose of measuring the Hubble constant well beyond the influence of local peculiar velocities. These NICMOS measurements mark the first time SBFs have been measured in galaxies out to redshifts of 10,000 \\kms, clearly demonstrating the advantages of measuring SBFs in the near-IR with excellent spatial resolution and low background. The calibration of the F160W SBF distance scale presented here was based on SBF measurements of galaxies in which Cepheid variable stars were detected in the same galaxy. Using a maximum-likelihood technique to account both for the influence of local mass concentrations on the velocity field and the non-Gaussian uncertainties on our SBF distance measurements yields a Hubble contsant of \\Ho$\\,{=}\\,76\\pm1.3$ \\kmsmpc\\ (1-$\\sigma$ statistical uncertainty) with an additional systematic uncertainty of 6 \\kmsmpc, primarily the result of uncertainty in the distance to the LMC. The small statistical uncertainty in \\Ho\\ is a result of the fact that our measurement is very insensitive to peculiar velocities, stellar population variations, extinction corrections, and photometric errors. Arbitrarily excluding all but the six most distant galaxies from the fit results in a 6\\% decrease in \\Ho, consistent with the hypothesis that the Local Group is located in an under-dense region of the universe."
    },
    "0011/astro-ph0011241_arXiv.txt": {
        "abstract": "A search for Gamma Ray Bursts in the GeV-TeV energy range has been performed by INCA, an air shower array working at 5200 m of altitude at the Chacaltaya Laboratory (Bolivia). The altitude of the detector and the use  of the ``single particle technique''  allows to lower the energy threshold up to few GeVs. No significant signals are observed during the occurrences of 125 GRBs detected by BATSE, and the obtained upper limits on the energy fluence in the interval 1-10$^{3}$ (1-10$^2$) GeV, range from 3.2 (8.6) 10$^{-5}$ to 2.6 (7.0) 10$^{-2}$ erg cm$^{-2}$ depending on the zenith angle of the events. These limits, thanks to the extreme altitude of INCA, are the lowest ever obtained in the sub-TeV energy region by a ground based experiment. ",
        "introduction": "Since their discovery in 1973 most of our knowledge on Gamma Ray Bursts (GRBs) has been confined to the KeV-MeV energy region; only during the last 3 years the observations extended to other wavebands, namely X, optical, infrared and radio, yielding a lot of exciting results and producing a break-through in our understanding of these misterious events (for a review see~\\cite{ref:gal} and~\\cite{ref:klo}). Now we know that GRBs are huge explosions occurring in far galaxies and are probably the most energetic events in the universe, exceeding even Supernovae. The cause of the release of such amount of energy is still unknown and what we observe is probably only the effect of relativistic shock waves produced by the explosion interacting each other or with the surrounding interstellar medium. Even if the ``central engine'' remains hidden the study of the radiation emitted in all wave bands by the interacting shock waves gives important informations on the acceleration mechanisms and on the ambient medium in which the shock  propagates (see~\\cite{ref:pir} for a theoretical review).  So far the radiation emitted by GRBs in the GeV-TeV band has been very poorly studied due to the extremely low fluxes. EGRET, the high energy detector aboard the satellite GRO, during its life detected only few very intense events containing some GeV photons~\\cite{ref:egr}. However all GRBs could contain a GeV energy component, so far not measured only because of the low flux. In fact, many models~\\cite{ref:mes,ref:vie,ref:bar,ref:der,ref:pao,ref:tot} predicts emissions in the GeV-TeV region. Upper limits in the 10-100 TeV region has been obtained by several air shower arrays ~\\cite{ref:eas,ref:cyg,ref:heg,ref:tib}. An interesting TeV candidate has been observed by the ground based Milagrito detector in coincidence with GRB970417~\\cite{ref:mil}.  The paucity of the flux is not the only problem to face when studing the high energy component of the GRB spectrum. The major problem (and unsolvable) is the absorption of gamma rays in the intergalactic space. GeV and TeV gamma rays interact with the infrared photons emitted by stars and dust producing electrons and positrons pairs. The flux of photons of energy $E$ decreases as $dN/dE = (dN_0/dE) e^{-\\tau(E,z)}$, where $z$ is the redshift of the source. The optical depth $\\tau$ increases with $E$ and $z$ but it is not easy to evaluate, due to the difficulty of measuring the infrared field in the far universe. According to a model~\\cite{ref:ste} the optical depth becomes equal to 1 for energies as low as $E \\sim$  40-70 (200-400) GeV when $z$=1(0.2).  So far about 15 redshifts of GRBs hosts have been measured: they range between 0.4 and 4.5, clustering around $z \\sim$ 1. This implies that even if GRBs emit TeV-PeV gamma-rays we unlikely could detect them, unless we are so lucky to observe a rare event occuring very close to us. As a consequence, to study the high energy component of GRBs we are forced to concentrate our efforts in the region of energy less than 1 TeV.  Ground based experiments can study high energy GRBs detecting the secondary particles of air showers generated by the interactions of gamma rays with the atmosphere nuclei. They can be divided in two major classes: Cerenkov Telescopes, detecting the Cerenkov photons emitted by the electrons and positrons traveling throught the atmosphere, and Air Shower Arrays, detecting the charge particles (mainly electrons and positrons) that reach the ground. Cerenkov telescopes are not suitable to detect transient events as GRBs, because of their small field of view (few squared degrees) and their limited duty cycle ($\\sim$ 10$\\%$). On the contrary, Air Shower Arrays have a larger field of view ($\\sim \\pi$ sr) and a duty cycle of $\\sim$ 100$\\%$. They usually work with energy thresholds of $\\sim$ 10-100 TeV, but in searching for transient events as GRBs, they can lower the threshold to few GeVs using the ``single particle technique'' and operating at very high altitude.  The single particle technique~\\cite{ref:ver,ref:eas} consists in counting all the particles giving a signal in the detector (not requiring any coincidence between particles as it is usually done to detect air showers). In this way it is possible to detect the lonely survivals of small showers produced by primaries of relatively low energy. Obviously with only one particle per shower it is not possible to reconstruct the arrival direction nor the energy of the primary and a GRBs is detectable only as a short duration excess over the single particle background, possibly in coincidence with a GRB satellite detection. The background is due to all secondary cosmic rays from all the sky above the horizon.  The signal from a GRB strongly increases with the altitude of the detector ~\\cite{ref:ver}. Fig.1 shows the mean number of particles reaching the ground generated by a gamma ray of different energies as a function of the altitude (the values refer to a zenith angle $\\theta$=30$^{\\circ}$). Going from 2000 to 5000 m, a 10 (100) GeV signal increases by a factor $f_s \\sim$100 (50); since the background increases only by a factor $f_b \\sim$ 4, the  sensitivity will improve by a factor $f = f_s/\\sqrt{f_b} \\sim$ 50 (25).  Following these considerations we decided to exploit the air shower array BASJE, operating since 1962 (with several upgrades since then) at the Chacaltaya Laboratory at 5200 m a.s.l., to develop the experiment INCA, to search for GeV GRBs using the single particle technique.  \\begin{figure}[htb] \\vfill \\begin{minipage}{.45\\linewidth} \\begin{center} \\mbox{\\epsfig{figure=np.eps,width=1.0\\linewidth}} \\end{center} \\vskip -0.5cm \\caption{Mean number of e$^{\\pm}$ reaching the ground generated by a gamma ray of different energies as a function of the altitude ($\\theta$=30$^{\\circ}$).} \\label{fi:np} \\end{minipage}\\hfill \\begin{minipage}{.45\\linewidth} \\begin{center} \\mbox{\\epsfig{figure=oscillo.eps,width=1.0\\linewidth}} \\end{center} \\vskip -0.5cm \\caption{Counting rate of 3 INCA scintillator detectors during $\\sim$ 10 days of measurement.} \\label{fi:osc} \\end{minipage}\\vfill \\end{figure} ",
        "conclusions": ""
    },
    "0011/astro-ph0011506_arXiv.txt": {
        "abstract": "\\rightskip = 0.0in plus 1em  We report on a search for long duration microlensing events towards the Large Magellanic Cloud.  We find none, and therefore put limits on the contribution of high mass objects to the Galactic dark matter. At 95\\% confidence level we exclude objects in the mass range $0.3 \\msun$ to $30.0 \\msun$ from contributing more than $4 \\ten{11} \\msun$ to the Galactic halo. Combined with earlier results, this means that objects with masses under $30 \\msun$ cannot make up the entire dark matter halo if the halo is of typical size. For a typical dark halo, objects with masses under $10 \\msun$ contribute less than 40\\% of the dark matter. ",
        "introduction": "\\label{sec-intro}  Recent results from the MACHO and EROS collaborations have ruled out Massive Compact Halo Objects (MACHOs) as the bulk the Galactic dark matter in the mass range $10^{-7}\\msun$ to a few solar masses \\cite{macho-lmc5,eros-2,macho-eros-spike}, thus eliminating the main candidate for baryonic dark matter in the Milky Way. However, there still remains a window between several solar masses and around one thousand solar masses \\cite{moore}, where black holes or other MACHOs, could make up the dark matter of the Milky Way. It was shown by Carr and Hawking (1974) that primordial black holes could have formed at very early stages in the Universe as a result of initial inhomogeneities, and recent work has focussed on the spike in the primordial black hole mass spectrum that could arise during the quark-hadron phase transition in the early Universe (e.g. Jedamzik 1997). As reviewed by Carr (1994), the density relative to critical density, $\\Omega$, in compact dark objects with masses of 0.01 -- 20 $\\msun$ must be less than 0.1 (increasing, however, to $\\simlt$1 for 60 -- 300 $\\msun$ objects),  limits determined by line-to-continuum microlensing effects in quasars (Dalcanton et al.~1997).  These limits still allow the Milky Way Halo to consist largely of such objects. For comparison, Galactic chemical enrichment arguments, assuming a standard initial stellar mass function and standard stellar evolution, limit the remnants of Population III stars to $\\Omega \\leq 0.001$ for remnants in the range 4 -- 200 $\\msun$ (Carr, Bond, \\& Arnett 1984).  However, there are no strong limits on black holes or non-topological soliton states arising from the early Universe or on relics from a non-standard very early generation of stars.  More recently, the microlensing surveys are setting the best limits on the fraction of dark objects in our own dark halo over a wide range of masses. The above collaborations search for MACHOs using the fact that when a compact dark object passes in front of a source star in a nearby dwarf galaxy, the source star suffers a temporary magnification due to gravitational microlensing \\cite{pac86}. The duration of this magnification is determined by a combination of the lens distance, velocity, and mass, and this degeneracy means that the lens mass cannot be uniquely determined for an individual lensing event unless other information is available.  However, for a given halo model an average over the lens density and velocity distributions can be made and an average duration can be estimated \\cite{griest91} \\begin{equation} \\that \\approx 130\\sqrt{m/\\msun} \\, \\rm days, \\label{eqt_hat} \\end{equation} where $\\that$ is the time for the source to cross the Einstein ring diameter.  The 13 -- 17 events discovered by the MACHO collaboration in the Large Magellanic Cloud (LMC) have durations ranging from about 30 days to about 130 days (Alcock et al.\\ 2000a), indicating lens masses in the $0.1\\msun$ to $1 \\msun$ range.  However, the number of events found, while larger than the expected background of 2--4 events from known stellar populations, is significantly less than the 60--80 events expected if the dark matter consisted entirely of objects in this mass range. Thus a most likely MACHO halo fraction of 20\\% was found, and a 100\\% MACHO halo was ruled out at 95\\% c.l. Earlier EROS \\cite{eros-ccd,macho-eros-spike} and MACHO \\cite{macho-spike,macho-eros-spike} searched for events with durations of less than 10 days and found none, limiting the MACHO fraction of dark matter to less than 25\\% in the range from $10^{-7}\\msun$ to $10^{-4}\\msun$. Finally, and most powerfully for the high mass end, the EROS collaboration \\cite{eros-2} did a combined analysis of all their microlensing surveys of the Magellanic Clouds, and set a 95\\% c.l. limit that objects in the $10^{-7}\\msun$ to $4 \\msun$ range do not constitute 100\\% of the dark halo, and objects less than $1 \\msun$ contribute less than 40\\% of the dark halo.  In this letter we improve on these limits in the high mass region by performing a search for events with durations longer than 150 days. We did not find any such events and therefore can improve the upper limit to around $30 \\msun$. ",
        "conclusions": "The limits given in this paper are  the strongest to date on compact halo objects with masses above 1 $\\msun$. In particular, black holes or other dark compact objects with masses less than $30 \\msun$ cannot make up the bulk of the dark matter.   We do note, however, that the present survey/analysis does not have much sensitivity to objects with masses greater $30 \\msun$. There is a large background of slowly varying variable stars which must be removed, and our main signal-to-noise cuts are not very good at distinguishing these from microlensing.  Thus we rely primarily on a long flat baseline, and on a direct cut on the fit event duration ($\\that<600$ days) to remove this background. Unfortunately, these cuts also limit our ability to detect long duration events coming from high mass lenses. We expect that analysis of the complete 8-year data set will go some way towards solving this problem and we expect to be able to push the current limit to higher masses, or to present long duration microlensing events when that dataset has been analyzed."
    },
    "0011/astro-ph0011395_arXiv.txt": {
        "abstract": "The relativistic precession model for quasi periodic oscillations, QPOs, in low mass X-ray binaries is reviewed. The behaviour of three simultaneous types of QPOs is well matched in terms of the fundamental frequencies for geodesic motion in the strong field of the accreting compact object for reasonable star masses and spin frequencies. The model ascribes the higher frequency kHz QPOs, the lower frequency kHz QPOs and the horizontal branch oscillations to the Keplerian, periastron precession and nodal precession frequencies of matter inhomogeneities orbiting close to the inner edge of the accretion disk. The remarkable correlation between the centroid frequency of QPOs in both neutron star and black hole candidate low mass X-ray binaries is very well fit by the model. QPOs from low mass X-ray binaries might provide an unprecedented laboratory to test general relativity in the strong field regime. ",
        "introduction": "Old accreting neutron stars, NSs, in low mass X-ray binaries, LMXRBs, display a complex variety of quasi-periodic oscillation, QPO, modes in their X-ray flux. The {\\it low frequency} QPOs ($\\sim 1-100$~Hz) that were discovered and studied from high luminosity Z-sources in the eighties are further classified into horizontal, normal and flaring branch oscillations (HBOs, NBOs and FBOs, respectively), depending on the simultaneous position occupied by a source in the X-ray colour-colour diagram (for a review see van der Klis 1995). The kHz QPOs ($\\sim 0.2$ to $\\sim 1.3$~kHz) that were revealed and investigated with RXTE in a number of NS LMXRBs (see van der Klis 1998, 1999, 2000 and references therein) involve timescales similar to the dynamical timescales close to the NS. A common phenomenon is the presence of a pair of kHz QPOs (centroid frequencies of $\\nu_1$ and $\\nu_2$) which drift in frequency while mantaining their frequency difference $\\Delta\\nu \\equiv \\nu_2 - \\nu_1 \\approx 250-360$~Hz roughly constant. Detailed studies showed that in four sources $\\Delta\\nu$ decreases significantly (by up to $\\sim 100$~Hz) as $\\nu_2$ increases; these are Sco~X-1 (van der Klis et al. 1997), 4U1608-52 (Mendez et al. 1998a,b), 4U1735-44 (Ford et al. 1998) and 4U1728-34 (Mendez \\& van der Klis 1999). Owing to poor statistics, a similar variation of $\\Delta\\nu$ in other sources would have remained undetected (Psaltis et al. 1998).  kHz QPOs show remarkably similar properties across NS LMXRBs of the Z and Atoll groups, the luminosity of which differs by a factor of $\\sim 10$ on average. During type I bursts from six Atoll sources, a nearly coherent signal at a frequency of $\\nu_{burst} \\sim 290 - 580$~Hz has also been detected (for a review see Strohmayer 2000). In a few cases $\\nu_{burst}$ is consistent, to within the errors, with the frequency separation of the kHz QPO pair $\\Delta\\nu$ or twice its value $2\\Delta\\nu$. Yet there are currently two sources (4U1636-53, Mendez et al. 1999, and 4U1728-34,  Mendez \\& van der Klis 1999) for which $\\nu_{burst}$ is significantly different from $\\Delta\\nu$ and its harmonics.  The presence of HBOs has been firmly established in both Atoll and Z-sources. Their frequency, $\\nu_{HBO}$ ($\\sim 15$ to $\\sim 60$~Hz) shows an approximately quadratic dependence ($\\sim \\nu_2^2$) on the higher kHz QPO frequency that is observed simultaneously in a number of sources. The frequency changes of the kHz QPOs and HBOs are positively correlated with the instantaneous accretion rate. Some evidence has also been found for an equivalent of the NBOs and FBOs of Z-sources (Wijnands et al. 1999; Psaltis, Belloni \\& van der Klis 1999).  A remarkable correlation between the centroid frequencies of QPOs (or peaked noise components) from LMXRBs has been recently discovered (Psaltis, Belloni \\& van der Klis 1999). This correlation extends over nearly 3 decades in frequency and encompasses both NS and black hole candidate, BHC, systems. The frequencies of these QPOs, despite their quasi-periodic nature, provide the most accurately measured observables of LMXRBs. A primary goal of any QPO model is therefore to explain the frequency range and dependence of the different QPO types of these sources. The basic features of the relativistic precession model, RPM, are reviewed here (Stella \\& Vietri 1998a, 1999; Stella, Vietri \\& Morsink 1999). In the RPM the QPO signals arise from the fundamental frequencies of the motion of matter in the vicinity of the NS. The corresponding orbits are supposed to be slightly eccentric and tilted. As in other models, the higher frequency kHz QPOs at $\\nu_2$ are produced by the $\\phi$-motion ({\\it i.e.} the Keplerian motion) of inhomogeneities orbiting the inner disk boundary, while the lower frequency QPO signal at $\\nu_1$ originates from their periastron precession, which is primarily determined by strong-field effects. The HBOs are due to the nodal precession in the orbits of the same inhomogeneities, an effect which is dominated by frame-dragging around fast-rotating collapsed stars. The RPM can be applied to BHCs as well. ",
        "conclusions": "\\subsection{Beat Frequency Models}  In the alternative scenario provided by beat frequency models, BFMs, disk inhomogeneities at the magnetospheric boundary ($r_m$) and the sonic point radius ($r_s$) are accreted at the beat frequency between the local Keplerian frequency, $\\nu_\\phi$, and the  NS spin frequency, $\\nu_s$, giving rise to the HBOs ($\\nu_{HBO}= \\nu_{\\phi}(r_m) - \\nu_s$) and the lower frequency kHz QPOs ($\\nu_1= \\nu_{\\phi}(r_s) - \\nu_s$), respectively (Alpar \\& Shaham 1985; Lamb et al. 1985; Miller et al. 1998). The higher frequency kHz QPOs are attributed to the Keplerian motion at the sonic point radius ($\\nu_2 = \\nu_\\phi(r_s)$). The frequency separation $\\Delta\\nu$ therefore, yields the NS spin frequency, $\\nu_s$. The narrow distribution of $\\nu_{s}$ inferred in this way ($\\sim 250-360$~Hz) is far from the mass shedding limit of any NS model and  considerably less extended than that of fast MSPs ($\\sim 600$~Hz). Accordingly the NSs of LMXRBs should be equilibrium rotators with a narrow frequency range despite their different average mass transfer rates, magnetic field strenghts and evolutionary histories. Moreover, if LMXRBs are the progenitors of MSPs, then BFMs would require a different evolutionary path leading to the formation of radio pulsars with $\\nu_s > 400$~Hz.  The fact that in some Atoll LMXRBs $\\nu_{burst}\\simeq \\Delta\\nu$ (or $\\simeq 2\\Delta\\nu$) is readily interpreted, because in BFMs $\\nu_s \\equiv \\Delta\\nu $. Yet, $\\Delta\\nu$ does vary and is significantly different from $\\nu_{burst}$ in several sources (see Sect.~1). This in contrast with the expectations of simple BFMs.  Attempts at fitting the PBV correlation within BFMs by using the range of spin frequencies inferred from $\\Delta\\nu$ fail to produce a power-law like behaviour over a sufficiently large range of frequencies (see van der Klis 1999). This is because both $\\nu_1$ and $\\nu_{HBO}$ result from the difference of a variable frequency ($\\nu_{\\phi}$ at $r_s$ and $r_m$, respectively) and a fixed frequency ($\\nu_s$). Moreover, BFMs are not applicable to BHCs, since the {\\it no hair theorem} excludes the possibility that an offset magnetic field or radiation beam can be stably anchored to the black hole, as required to produce the beating with the disk Keplerian frequency.  \\subsection{The Relativistic Precession Model}  The RPM naturally explains the frequency range and dependence of the kHz QPOs and HBOs in NS LMXRBs, as well as the PBV correlation, which involves both NS and BHC systems. The model has a minimum number of free parameters. Predictions that can be tested through future analyses and/or observations include:  (a) The frequency difference $\\Delta\\nu=\\nu_r$ is expected to decrease also for low values of $\\nu_2=\\nu_\\phi$ (see Eq.~4 and Fig.~1). Moreover if the highest $\\nu_2=\\nu_\\phi$ frequencies do originate from nearly circular orbits (see Fig.~1A), then $\\Delta\\nu$ should quickly decrease as $\\nu_2$ increases further and the orbital radius approaches the marginally stable orbit.  (b) $\\nu_2=\\nu_\\phi$ is expected to scale as $\\nu_1^{3/5} = \\nu_{per}^{3/5}$. Extending the $\\nu_2$ vs. $\\nu_1$ correlation in NS systems toward lower frequencies and detecting the signal at $\\nu_2$ in BHC systems would provide important new tests.   In the RPM the QPO signals are produced at $ r/M \\sim 100\\ m^{-2/5}\\nu_{per}^{-2/5}$, which, for individual sources, must decrease for increasing mass accretion rates (as $\\dot M$ is positively correlated with {\\it e.g.} $\\nu_2$). The inferred radii range from close to the marginally stable ($r/M \\sim 6$) to $r/M \\sim 30 $ over the frequency span covered by the PBV correlation. Many NS and BHC LMXRBs display two component  X-ray spectra consisting of a soft thermal component, usually interpreted in terms of emission from an optically thick accretion disk, and a harder, often power-law like component, likely due to a hot inner disk region. The QPOs might originate at the transition radius between the optically thick disk and the hot inner region as a result of occultation or modulated emission by inhomogeneities (Stella, Vietri \\& Morsink 1999; see also Di Matteo \\& Psaltis 1999).   Mechanisms that can induce a finite (though small) eccentricity and tilt in the motion of matter in the innermost disk regions are currently under investigation. In the case of NSs, some kind of resonance between the star spin and the motion of matter inhomogeneities (see {\\it e.g.} Vietri \\& Stella 1998) might be responsible for the close commensurability of $\\nu_{burst}$ and $\\Delta\\nu$ (or $2\\Delta\\nu$) observed in a few sources. These issues will be addressed elsewhere.  A simple test-particle approximation has been adopted so far within the RPM. Much needed hydrodynamical calculations are still in their infancy; among other things these are hampered by uncertainties concerning the physics of the innermost disk regions. Yet we note that in the hydrodynamical approach explored by Psaltis \\& Norman (2000), the test particle frequencies (the same as in the RPM plus an additional frequency at $\\nu_\\phi+\\nu_r$) are selected by the response of an annulus in the disk, when this is subject to a wide-band input noise.  If confirmed, the RPM will provide an unprecedented opportunity to measure GR effects in the strong field regime, such as the periastron precession in the vicinity of the marginally stable orbit and the radial dependence of the Lense-Thirring nodal precession frequency. In principle, accurately measured kHz QPO and HBO frequencies would yield crucial information on the compact object such as its mass and angular momentum ({\\it e.g.} by solving Eqs. 1-3 for $m$, $a/M$ and $r$). Should suitable, additional observables be found, it might become possible to obtain a self-consistency check of the RPM, together with tests of GR in the strong field regime."
    },
    "0011/astro-ph0011489_arXiv.txt": {
        "abstract": "We have detected coherent   oscillations, at multiple frequencies,  in the  line and  continuum emission of  the  eclipsing dwarf  nova V2051 Ophiuchi using  the   10m Keck II    telescope.  Our own  novel   data acquisition system allowed us to obtain very fast spectroscopy using a continuous readout of  the CCD on the LRIS  spectrograph.  This is the first time  that dwarf nova  oscillations are detected and resolved in the emission lines.  The accretion  disc  is highly asymmetric with  a stronger contribution  from the blue-shifted side  of the  disc during our observations.  The disc extends from  close to the white dwarf out to the outer regions of the primary Roche lobe.  Continuum oscillations at 56.12s and its first harmonic at 28.06 s are most likely to originate on the surface of a spinning white dwarf with the fundamental period corresponding to the spin period.  Balmer and Helium emission lines oscillate with  a period of 29.77s  at a mean amplitude of  1.9\\%.  The line kinematics  as  well as  the eclipse  constraints indicate  an   origin   in  the   accretion   disc  at a   radius   of $12\\pm2R_{wd}$.  The    amplitude of  the  emission  line  oscillation modulates  (0-4\\%) at  a period  of  488s, corresponding to the Kepler period at R=$12R_{wd}$. This modulation is due to the beating between the white dwarf spin and the orbital motion in the disc.  The  observed  emission line oscillations  cannot  be  explained  by a truncated disc as in the intermediate polars. The observations suggest a   non-axisymmetric bulge in the    disc,  orbiting at $12R_{wd}$,  is required.     The close correspondence  between   the  location of the oscillations  and the  circularisation radius  of the system  suggests that stream overflow effects may be of relevance. ",
        "introduction": "Rapid variability is a     tell-tale sign of accretion.   The    first coherent oscillation in a cataclysmic  variable binary was reported by Walker  as far back as 1956  (Walker  1956\\nocite{w56}).  The strictly periodic 71s    brightness oscillations he   found in  DQ Her  are now understood as being produced  by a spinning  magnetic white dwarf that is being fed by an accretion disc, and accretes the material along its magnetic field lines. Since the dynamical  time scales of matter close to the white dwarf  are of the  order of tens of  seconds, it  was not until the development of high speed photon-counting photometers in the '70s  that periodic signals   were discovered in  other  systems.  The dwarf nova   oscillations (DNOs) are periodic   brightness modulations discovered in dwarf nova systems during outburst (Warner \\& Robinson  1972\\nocite{wr72}).    The   short time   scales  (2-40s) indicate  an   origin close to   the   white dwarf. In  some  systems, different   periods  were found   at  different times, and  the period stability is considerably lower  than the coherent oscillations  in DQ Her.   See Warner   (1995)\\nocite{warner}  for  review.  Whereas   the amplitudes  in  the optical are   usually  so low ($<  1\\%$)  that the oscillations are only revealed after period analysis, Marsh \\& Horne  (1998)\\nocite{mh98}  detected coherent  oscillations at  two periods in the HST light curves of OY  Car with an  amplitude of up to 20\\% towards the end of a super-outburst.  Reminiscent  of the  DNOs are   the flood of  QPOs   reported in X-ray observations of accreting neutron    star binaries (see van   der Klis 2000\\nocite{klis} for  review).  The properties of  these oscillations are in many ways very similar to those of DNOs, albeit on shorter time scales since these systems contain a  much smaller compact object.  It shows that moderately  coherent oscillations  appear  to be  a generic feature of accretion onto compact objects.  V2051 Ophiuchi is a short period eclipsing dwarf nova system belonging to the  SU UMa sub-class.  Its status  as a  dwarf nova was contested (e.g.    Warner  \\&    O'Donoghue    1987\\nocite{wo87})   until    two super-outbursts were observed  in 1998  and  1999 by amateur  variable star networks.  Spectroscopy reveals double-peaked emission lines that exhibit a  rotational disturbance effect during eclipse, corresponding to      a  pro-grade   rotating  accretion  disc       (Cook  \\& Brunt 1983\\nocite{cb83}, Watts  et  al.  1986\\nocite{watts86}).  Baptista et al.  (1998)\\nocite{bap98} present a  photometric model for the  binary based on HST   and  ground based   observations.  Due  to  the unusual low-state the system was   in during their observations,  the  orbital phases when bright spot and white dwarf are eclipsed by the mass donor star could be directly measured.  Warner \\& Cropper (1983)\\nocite{wc83} concluded that most of the rapid variability (flickering) was produced in the inner disc regions rather than   the  bright spot   or white  dwarf.   The  only quasi  periodic oscillation  in   this system  was  reported  by Warner  \\& O'Donoghue (1987), who  detected  a very  short   lived  42s  oscillation  during outburst.   It was  not  classified  as  a  DNO because   of  its low coherence,  and the  absence of  a   corresponding spike in the  power spectrum at that period.  We present high time resolution spectroscopy of the dwarf  nova V2051 Oph shortly after  a normal outburst, obtained using our own data collection hardware equipped on  Keck II, and report on the discovery of emission line oscillations at the level of several percent.  In the next  Section we will  describe the data acquisition system and novel reduction steps that were used. Section 3 covers the analysis of the data and discusses the properties of the oscillations.  We discuss likely models in Section 4 and summarise our conclusions in Section 5. ",
        "conclusions": "We detected  coherent continuum and emission  line oscillations in the dwarf nova V2051 Oph on decline from a normal outburst. Accretion disc emission extends from  very close to the  white dwarf out to the outer parts of the  primary Roche lobe.  The disc  emission lines  display a persistent blue to red asymmetry, with the blue peak being stronger in all the Helium  lines as  well as  the Balmer lines,  except H$\\alpha$ which has a stronger red peak. The eclipse light curves are also highly asymmetric  which  suggests  that the blue  side  of  the disc makes a larger contribution to the emission compared to the red-shifted side.    The  continuum oscillations are most likely to originate on the surface of a spinning white dwarf with spin period 56.12s and  temperature around 15000K.  The amplitude of the oscillation   in  the continuum  varies  between 0  and   4\\%, and disappears  during   white dwarf eclipse.   The   Balmer  and Helium I emission  lines oscillate strongly at  a  period  of 29.8s. The line kinematics as well as  the eclipse constraints indicate these  to come from a non-axisymmetric bulge in the disc at   a  radius of $12\\pm2R_{wd}$.   The corresponding Kepler orbit has a  period of 488s, and corresponds well with the  observed amplitude variations (0-4\\%)  on this  period.  The oscillating line emission is observed to go from maximum blue-shift to maximum red-shift every 29.8s, but not vice versa. This, together with the  regular  modulation   of  the  oscillation  amplitudes   can be explained by intervening  gas  above the orbital plane,  close  to the white dwarf. This is supported by an uneclipsed component in both the continuum and lines. In the Balmer lines this component is centered on zero velocity and has a FWHM of 1000 km~s$^{-1}$. Alternatively, the emission is beamed away from the white dwarf.  The close correspondence between the  location of the oscillations and the circularisation radius of the system as well as the disc asymmetry may indicate   the relevance of  stream  overflow  to the  presence of vertically extended bulges in the  disc. More observations of this kind as well as more detailed simulations of such a  scenario are required to confirm its feasibility in the light of producing DNOs. Predictions of our interpretation are  the persistent  presence of 56.12s/28.06s oscillations from  the white dwarf.  This  period should be detected  at other  epochs, most  likely in  the  UV.  The UV  also provides opportunities to measure the spin of the white dwarf directly using the rotational broadening  of white dwarf absorption lines (Sion 1999\\nocite{sion}). With a spin period of 56.12s, the expected surface velocity of the spinning white dwarf works out to be: \\[ v_{wd} = \\Omega \\times R_{wd} = 2\\pi/P_2 \\times R_{wd} = 810~ km~s^{-1} \\] Secondly,   we   expect  the  period  of  the    emission line oscillations to  change across the outburst  cycle as the location and extent of  stream  overflow and re-impact  will  depend on  the radial density distribution of the disc. For example, if we wish to interpret the reported 42.2s oscillation (Warner \\& O'Donoghue 1987) in this scenario, it can be identified as the beat between the white dwarf spin and a Keplerian orbit at $\\sim 7 R_{wd}$.  The data presented   here demonstrate the   potential  of high  speed spectroscopy of  CVs  using large   aperture telescopes. We   detected emission line oscillations  in a dwarf nova for  the first time,  and located their origin in the accretion disc.  A better understanding of the  elusive DNOs would  benefit from a high  time resolution study of (eclipsing) dwarf novae across their outburst cycle."
    },
    "0011/astro-ph0011440_arXiv.txt": {
        "abstract": "We present a distance to \\object{NGC\\,4527}, the host galaxy of the type Ia \\object{SN\\,1991T}, measured by surface brightness fluctuations. This supernova has been labelled ``peculiar'' both on the grounds of its spectroscopic behaviour and its apparent overluminosity with respect to other supernovae. The distance modulus to \\object{NGC\\,4527} and thus to \\object{SN\\,1991T} is $30.26\\pm0.09$. This relatively short distance largely removes the discrepancy with other Ia supernovae having similar light-curve characteristics and also removes the motivation for interpreting \\object{SN\\,1991T} as a super-Chandrasekhar explosion.  However, the reddening uncertainty results in significant uncertainty of the absolute magnitudes.  We show further that, although the intrinsic absolute magnitudes among Ia supernovae are quite different, their brightness, corrected for decline-rate and colour, shows a high degree of uniformity, including underluminous events like \\object{SN\\,1991bg} and \\object{SN\\,1997cn}. In particular, the present data do not support a significant difference between corrected SNe Ia luminosities in late-type and early-type host galaxies. ",
        "introduction": "A precondition for the successful use of Ia supernovae at high redshifts for constraining cosmological models is the understanding of low-redshift supernovae, in particular their luminosities and their dependence on light-curve parameters such as the decline-rate \\cite{phillips:93}. Besides the decline rate, it is also the colour at maximum, for which one has to correct \\cite{tripp:98,drenkhahn:99}. The Cal\\'an-Tololo sample \\cite{hamuy:96} and the Harvard sample \\cite{riess:99} demonstrated that most Ia supernovae, after appropriate corrections, exhibit a surprisingly high degree of uniformity in their maximum brightness.  A few supernovae, however, seem to deviate in the sense that they are either too bright or too dim. Since at high redshifts, selection effects towards the bright end of the supernovae luminosity distribution are probably present, the apparently overluminous supernovae are of particular interest.  Obviously, the distance to the host galaxy is the key parameter.  For example, in the case of \\object{SN\\,1994D} in \\object{NGC\\,4526}, it could be shown that this apparently overluminous Ia event took place at a shorter distance than initially assumed, and thus its luminosity is consistent with other type-Ia SNe \\cite{drenkhahn:99}.  Another more extreme example of overluminosity is \\object{SN\\,1991T} in \\object{NGC\\,4527}. Besides its allegedly high luminosity, it also exhibited spectroscopic peculiarities.  Fisher et~al. \\cite*{fisher:99} discuss the present knowledge about \\object{SN\\,1991T} and also offer an interpretation of these abnormal spectral features.  Their analysis suggested that \\object{SN\\,1991T} was a super-Chandrasekhar explosion following the merging of two white dwarfs.  Although there is general agreement that \\object{NGC\\,4527} is located on the near side of the Virgo cluster, there is no individual distance determination besides a Tully-Fisher distance, where it is known that large deviations can occur. Fisher et~al. \\cite*{fisher:99} assumed its distance to be equal to the distance of two neighboring spiral galaxies, which also have been hosts for Ia SNe, \\object{NGC\\,4496} (\\object{SN\\,1960F}), and \\object{NGC\\,4536} (\\object{SN\\,1981B}).  The aim of our paper is to present an individual distance to \\object{SN\\,1991T}, based on surface brightness fluctuations in the bulge of \\object{NGC\\,4527}, which is more accurate than what is available in the literature.  Before we present our method of distance determination and its result, we recall briefly what is known about the distance of this galaxy group and the luminosities of the respective Ia SNe.   \\subsection{What is known about the distances of \\object{NGC\\,4527} and \\object{NGC\\,4536}?}  \\object{NGC\\,4536}, the host of \\object{SN\\,1981B}, has a more or less well established distance, based on the PL-relation for Cepheids observed by HST \\cite{saha:96a,saha:96b}. However, the Cepheid data are to a minor degree subject to interpretational manipulation, depending on details of the data reduction and also of the selection of Cepheids used for the distance determination.  Saha et~al. \\cite*{saha:96b}, for instance, quote a distance modulus to \\object{NGC\\,4536} of $31.07$\\,mag, while Gibson et~al. \\cite*{gibson:00} give $30.95\\pm0.23$\\,mag.  Tully-Fisher (TF) distances have been measured both for \\object{NGC\\,4527} and for \\object{NGC\\,4536}. If we use the new TF calibrations given by Sakai et al.  \\cite*{sakai:00} and the TF data given by Pierce \\cite*{pierce:94}, we then get distance moduli (averaged over $B$, $R$, $I$) of $30.78$ and $30.80$ for \\object{NGC\\,4527} and \\object{NGC\\,4536}, respectively. Pierce \\cite*{pierce:94} himself quotes $30.6$ and $30.5$, respectively.  The differences are consistent with the dispersion of the TF relation which is about $0.3$--$0.4$\\,mag.  In this sense, the TF data are in agreement with the Cepheid distance of \\object{NGC\\,4536}. Moreover, the TF distances cannot support any distance difference between \\object{NGC\\,4527} and \\object{NGC\\,4536} and it appears reasonable to assume a distance modulus of $30.95\\pm0.23$\\,mag for \\object{NGC\\,4527} as well. ",
        "conclusions": "By applying the method of surface brightness fluctuations to the bulge of \\object{NGC\\,4527}, we measured an accurate distance to this galaxy, which hosted \\object{SN\\,1991T}. Our result for the distance modulus is $30.26\\pm0.09$.  With this distance, the absolute maximum magnitudes of \\object{SN\\,1991T}, corrected for decline rate and colour, in $B$, $V$, $I$ are $-19.18\\pm0.52$, $-19.16\\pm0.39$, and $-18.95\\pm0.18$, respectively (adopting a reddening of $E(B-V)=0.2\\pm0.1$.  Although the uncertainty is large, mainly caused by the uncertain reddening, there is no reason to conclude that \\object{SN\\,1991T} has a peculiarly high luminosity, as the comparison with other SNe shows. Thus the model of a super-Chandrasekhar explosion advanced by Fisher et~al.  \\cite*{fisher:99} can not be motivated by the new distance.  Since also underluminous events like \\object{SN\\,1991bg} and \\object{SN\\,1997cn} are consistent with the corrected absolute magnitudes of the other SNe, there is now no single SN known, which deviates strikingly, in accordance with the homogeneous samples of the Cal\\'an-Tololo and the Harvard sample. Since there is evidence that SNe Ia in late-type and early-type galaxies stem from different parent populations, (see Leibundgut \\cite*{leibundgut:00} for a general review on SNe Ia) this may indicate that there is no evolutionary effect as far as the corrected absolute magnitudes are concerned. This is encouraging for the use of SNe Ia in cosmological contexts."
    },
    "0011/astro-ph0011326_arXiv.txt": {
        "abstract": "Using the observed properties of our solar system, in particular the isotopic compositions of meteorites and the regularity of the planetary orbits, we constrain the star formation environment of the Sun within the scenario of (external) radioactive enrichment by a massive star. This calculation yields a probability distribution for the number of stars in the solar birth aggregate. The Sun is most likely to have formed within a stellar group containing $N = \\nexp \\approx 2000$ $\\pm$ 1100 members.  The {\\it a priori} probability of a star forming in this type of environment is ${\\cal P} \\approx$ 0.0085, i.e., only about 1 out of 120 solar systems are expected to form under similar conditions. We discuss additional implications of this scenario, including possible effects from the radiation fields provided by the putative cluster environment and dynamical disruption of the Kuiper Belt. The constraints of this paper place tight restrictions on the properties of the solar birth aggregate for the scenario of external enrichment by a massive star; alternately, these tight constraints slightly favor a self-enrichment scenario for the short-lived radioactive species. ",
        "introduction": "A substantial fraction of star formation in our galaxy takes place in groups or clusters, i.e., associations containing many young solar systems.  As a consequence, our own solar system may have formed within such a crowded environment. This hypothesis is bolstered by observations of meteorites, which indicate that unexpectedly large quantities of short-lived radioactive nuclei were present at the epoch of planet formation. One traditional explanation for this set of abundance anomalies is that the solar nebula was enriched in radioactive species by a nearby massive star (see, e.g., the recent review of Goswami and Vanhala 2000; Cameron et al. 1995; Cameron 1978).  Enrichment is usually envisioned to occur through a supernova explosion, although Wolf-Rayet winds have also been invoked. In any case, this scenario requires the presence of a nearby massive star, which in turn implies that our solar system formed within a reasonably large stellar group.  For completeness, we note that thermally pulsing asymptotic giant branch stars have also been suggested as an enrichment source (see Brusso, Gallino, and Wasserburg 1999), but the probability that such a star is associated with a molecular cloud is relatively low (Kastner and Myers 1994).  We also know that our solar system could not have formed within a cluster containing too many stellar members. In a sufficiently crowded environment, the solar system would be disrupted by gravitational scattering effects from passing stars, binary systems, and other solar systems. The observed orbital elements of the outer planets exhibit low eccentricities and are (almost) confined to the same orbital plane. This relatively well-ordered configuration places modestly tight restrictions on the characteristics of the solar birth aggregate.  Some previous work has addressed this issue.  The inclinations of the orbits of Uranus and Neptune are sensitive to gravitational perturbations from passing stars in the solar birth cluster; this effect implies a bound (Gaidos 1995; see also Tremaine 1991) on the product of the stellar number density and the residence time in the cluster: $\\int n dt < 3 \\times 10^{4}$ pc$^{-3}$ Myr. In this present work, we constrain the size of the solar birth aggregate from both directions: It must be large enough to provide a sufficiently massive star for enrichment and small enough to allow for well-ordered planetary orbits.  We must keep in mind, however, another possible way to explain the inferred abundances of radioactive species: The solar nebula can be self-enriched through energetic processes in its early formative stages (e.g., Lee et al. 1998). This latter scenario allows for the solar system to form in relative isolation.  Recent evidence (McKeegan et al. 2000) provides support for the self-enrichment scenario. This work indicates the presence of $^{10}$Be, a short-lived radioactive species that cannot be produced in supernovae and hence must be produced by nuclear reactions from energetic particles (as would be expected in a picture of self-enrichment).  An important issue for solar system formation is to decide between these two alternate enrichment scenarios. As we show here, the external enrichment scenario places strong constraints on the birth environment of the solar system. These tight constraints, in turn, imply a low probability of success for the external enrichment scenario and thus suggest that self-enrichment may be more viable.  In this paper, we explore the compromise necessary to have the solar system formed within a stellar group large enough to produce a nearby massive star and, simultaneously, small enough to allow the planetary orbits to remain relatively unperturbed. We first determine the probability of a stellar cluster containing a sufficiently massive star as a function of the number $N$ of cluster members (\\S 2.1). We then calculate the cross sections for passing stars to disrupt the orbits of the outer planets in the solar system (\\S 2.2), and use the results to estimate the probability of disruption as a function of cluster size $N$ (\\S 2.3).  Folding these results together, we find the probability distribution for the size $N$ of the solar system's birth aggregate (\\S 2.4), the corresponding expectation value $\\langle N \\rangle$ for the cluster size, and the {\\it a priori} probability of a star forming in such an environment (\\S 2.5). Next, we use these results to reconstruct the expected ultraviolet radiation field impinging upon the solar nebula during the epoch of planet formation (\\S 2.6). Finally, we calculate the cross sections for the scattering of Kuiper Belt objects through gravitational interactions with passing stars (\\S 2.7) and thereby place further constraints on the birth aggregate. We conclude (\\S 3) with a summary and discussion of our results. ",
        "conclusions": "In this paper, we have explored the consequences of the solar system being formed within a group environment that is large enough to contain a massive star that enriches the early solar system in radioactive species and is also sufficiently diffuse to allow the planetary orbits to remain unperturbed. In particular, we have obtained the following results:  [1] We have calculated the cross sections for the outer planets in our solar system to experience orbital changes due to scattering interactions with binary systems in a cluster environment. We find the cross sections for eccentricity increases (see Table 1) and for increases in the relative inclination angles of the planetary orbits (see Table 2). The cross section for the scattering events to increase either the eccentricities or the inclination angles beyond the currently observed values is $\\cross \\approx$ (400 AU)$^2$. The cross section for planetary ejection and/or capture is somewhat lower, about $\\cross \\approx$ (130 AU)$^2$. However, all of these cross sections are substantially larger than the area subtended by the solar system, {\\sl Area} $\\approx \\pi$ (30 AU)$^2$.  [2] We have estimated the probability distribution for the number $N$ of stars in the birth aggregate for our solar system (see Figure 1). Using the coupled constraints that the group was large enough to contain a massive star (to enrich the solar system in radioactive elements) and small enough so that the outer planetary orbits are not severely disrupted, we find that $N = \\nexp \\approx 2000 \\pm 1100$. The expectation value $\\nexp$ varies relatively slowly with the parameters of the problem and has the analytic solution given by equation [\\ref{eq:ngeneral}].  [3] The {\\it a priori} probability for a star being born in the type of environment required for the external enrichment scenario for our solar system (i.e., subject to the probability distribution depicted in Figure 1) is ${\\cal P} \\approx 0.0085$.  The odds of the solar system forming in this type of environment are thus about 1 in 120. This result can be readily generalized to accommodate other choices of parameters (see eq. [\\ref{eq:pseries}] and Figure 2).  [4] The time scales for the cluster environment obey a particular ordering (see equation [\\ref{eq:timeorder}]): The crossing time $t_{\\rm cross}$ of the cluster is much shorter than the time $t_P$ required to form giant planets, so the solar system randomly samples the cluster environment over that epoch. The dynamical relaxation time $t_R$ is somewhat longer than $t_P$, and the total cluster lifetime $\\tcl \\gg t_P$, so the cluster does not change its structure appreciably while the planets form.  [5] Within the scope of this external enrichment scenario, we have reconstructed the radiation field provided by the birth aggregate of the solar system.  The early solar nebula receives about 40 times more ionizing UV photons from the background cluster environment than from the early Sun. Only the inner portion of the nebula, at radii $r < 5$ AU, has its ionizing UV flux dominated by the Sun.  This intense flux of UV radiation can severely ablate the early solar nebula and greatly compromise giant planet formation.  [6] Objects forming in the Kuiper Belt will also be scattered due to interactions in the birth cluster. We have calculated the scattering cross sections for KBOs on initially circular orbits with radii in the range 30 AU $\\le a \\le$ 70 AU (Table 3). The cross sections for prompt removal are $\\cross \\approx$ (350 AU)$^2$ over the outer part of this radial range. If we include the required survival of the Kuiper Belt in the probability analysis, we obtain slightly tighter constraints on the solar birth aggregate: $\\nexp$ = 1520 $\\pm$ 827, $\\cal P$ = 0.0067, and {\\it a priori} odds of one part in 150.  Some authors have suggested that the formation of our solar system must be triggered by the same supernova that is postulated to provide the radioactive enrichment (e.g., Boss and Foster 1998). We stress here that external enrichment does not necessarily imply a triggered collapse.  The time scale for cluster formation is relatively short (a few Myr; e.g., Elmegreen 2000) and the time scale for the collapse of an individual star forming site is much shorter (about 10$^5$ yr; e.g., Myers and Fuller 1993, Adams and Fatuzzo 1996). Thus, the solar system could be formed within a cluster and yet be formed through a spontaneous (un-triggered) collapse. Another motivation for the collapse being triggered is that the ambipolar diffusion time scale is generally long (about $10^7$ years), too long for the survival of the necessary radioactive nuclei. However, observational evidence shows that the time scale for molecular cloud cores to shed their magnetic support and begin dynamic collapse is much shorter (about 1 Myr) for all cores (e.g., Jijina et al. 1999; Myers and Lazarian 1998); the ambipolar diffusion time scale is thus not an insurmountable obstacle for star formation in clusters.  The results of this paper have important ramifications for the ongoing debate concerning radioactive enrichment of the early solar system. The meteoritic data strongly indicate that such enrichment took place. However, both the external scenario of enrichment by a massive star and the internal scenario of self-enrichment have some difficulty reproducing all of the short-lived radioactive species (e.g., see Goswami and Vanhala 2000; Lee et al. 1998; and references therein). Although the results are not definitive, this paper tends to favor the self-enrichment scenario for two reasons: (1) A star forming environment that is simultaneously large enough to provide external enrichment and diffuse enough to not disrupt the planetary orbits is {\\it a priori} an unlikely event (at the level of 1 part in 120). (2) If the solar system formed in the large birth aggregate required for external enrichment, the corresponding radiation fields are likely to have compromised giant planet formation.  Nevertheless, the external enrichment scenario is not conclusively ruled out. Events with long odds (100:1) do indeed sometimes happen. If one adopts the external enrichment scenario, a consistent solution exists and the results of this paper place tight constraints on the birth environment of our solar system: The birth cluster contained $N \\approx 2000$ other stars and subjected the early solar nebula to an intense UV radiation field. Scattering interactions in the birth cluster would then be responsible (at least in part) for the observed nonzero (but small) inclination angles and eccentricities of the giant planet orbits.  Furthermore, objects in the Kuiper Belt would be dynamically excited by these scattering interactions. This highly disruptive environment must be accounted for in a complete description of solar system formation.  \\bigskip  During the preparation of this paper, we have greatly benefited from discussions with Gary Bernstein, Pat Cassen, Gus Evrard, Dave Hollenbach, Nathan Schwadron, Frank Shu, and Kevin Zahnle. This work was supported by NASA Grant No. 5-2869 and the University of Michigan.  \\newpage  \\bigskip \\centerline{} \\bigskip  \\centerline{\\bf Table 1} \\medskip \\centerline{\\bf Cross Sections for Planet Scattering: Eccentricity Increase} \\bigskip \\centerline{[all cross sections in units of (AU)$^2$]} \\bigskip  \\begin{center} \\begin{tabular}{ccccc} \\hline \\hline \\hline $\\epsilon$ & Jupiter & Saturn & Uranus & Neptune \\\\ \\hline \\hline 0.05& 54300 $\\pm$ 707 & 67700 $\\pm$ 775&126000 $\\pm$ 1030 & 167000 $\\pm$ 1150\\\\ 0.10& 40700 $\\pm$ 611 & 55000 $\\pm$ 700&106000 $\\pm$  958 & 143000 $\\pm$ 1080\\\\ 0.20& 29200 $\\pm$ 521 & 42900 $\\pm$ 621& 83300 $\\pm$  860 & 113000 $\\pm$ 983\\\\ 0.30& 22800 $\\pm$ 458 & 36700 $\\pm$ 576& 69500 $\\pm$  794 & 95600 $\\pm$  915\\\\ 0.40& 18600 $\\pm$ 410 & 32400 $\\pm$ 544& 59300 $\\pm$  738 & 82900 $\\pm$  861\\\\ 0.50& 15500 $\\pm$ 374 & 28800 $\\pm$ 517& 52300 $\\pm$  697 & 73300 $\\pm$  816\\\\ 0.60& 13700 $\\pm$ 351 & 25800 $\\pm$ 491& 46600 $\\pm$  661 & 64900 $\\pm$  774\\\\ 0.70& 11900 $\\pm$ 327 & 22700 $\\pm$ 462& 41500 $\\pm$  626 & 58700 $\\pm$  740\\\\ 0.80& 10500 $\\pm$ 306 & 19800 $\\pm$ 432& 36700 $\\pm$  592 & 53200 $\\pm$  710\\\\ 0.90&  9120 $\\pm$ 286 & 17500 $\\pm$ 407& 32300 $\\pm$  558 & 47700 $\\pm$  678\\\\ 0.95&  8520 $\\pm$ 276 & 16400 $\\pm$ 393& 30400 $\\pm$  542 & 45100 $\\pm$  662\\\\ 1.00&  7970 $\\pm$ 267 & 15300 $\\pm$ 381& 28400 $\\pm$  526 & 41900 $\\pm$  640\\\\ \\hline \\hline escape& 7290 $\\pm$ 255 &14000 $\\pm$ 365 &25000 $\\pm$ 492 &35400 $\\pm$ 584\\\\ capture& 684 $\\pm$ 81.4 &1300 $\\pm$ 115 & 3320 $\\pm$ 194 & 6640 $\\pm$ 280\\\\ \\hline \\hline \\hline \\end{tabular} \\end{center}  \\newpage \\bigskip \\centerline{} \\bigskip  \\centerline{\\bf Table 2} \\medskip \\centerline{\\bf Cross Sections for Planet Scattering: Inclination Increase} \\bigskip \\centerline{[all cross sections in units of (AU)$^2$]} \\bigskip  \\begin{center} \\begin{tabular}{cccc} \\hline \\hline \\hline $\\Delta \\theta_i$ & angle increase & total \\\\ \\hline \\hline 0.06 & 140900 $\\pm$1100 & 158000 $\\pm$1130\\\\ 0.10 & 114600 $\\pm$1010 & 132000 $\\pm$1050\\\\ 0.20 &  85000 $\\pm$ 880 & 102000 $\\pm$ 943\\\\ 0.30 &  71800 $\\pm$ 815 &  89000 $\\pm$ 888\\\\ 0.40 &  63600 $\\pm$ 770 &  80800 $\\pm$ 850\\\\ 0.50 &  57600 $\\pm$ 736 &  74800 $\\pm$ 822\\\\ 0.60 &  53100 $\\pm$ 707 &  70300 $\\pm$ 798\\\\ 0.70 &  49500 $\\pm$ 684 &  66700 $\\pm$ 780\\\\ 0.80 &  46400 $\\pm$ 663 &  63600 $\\pm$ 763\\\\ 0.90 &  43900 $\\pm$ 647 &  61000 $\\pm$ 750\\\\ 1.00 &  41400 $\\pm$ 629 &  58500 $\\pm$ 736\\\\ 1.10 &  39500 $\\pm$ 616 &  56600 $\\pm$ 726\\\\ 1.20 &  37700 $\\pm$ 602 &  54800 $\\pm$ 715\\\\ 1.30 &  35700 $\\pm$ 585 &  52800 $\\pm$ 702\\\\ 1.40 &  33700 $\\pm$ 570 &  50900 $\\pm$ 690\\\\ 1.50 &  32400 $\\pm$ 559 &  49500 $\\pm$ 682\\\\ 1.60 &  30900 $\\pm$ 546 &  48000 $\\pm$ 672\\\\ 1.70 &  29400 $\\pm$ 534 &  46500 $\\pm$ 663\\\\ 1.80 &  27700 $\\pm$ 518 &  44900 $\\pm$ 652\\\\ 2.00 &  24600 $\\pm$ 488 &  41700 $\\pm$ 630\\\\ 2.20 &  21400 $\\pm$ 459 &  38500 $\\pm$ 609\\\\ 2.40 &  17100 $\\pm$ 411 &  34300 $\\pm$ 576\\\\ 2.60 &  12300 $\\pm$ 349 &  29400 $\\pm$ 537\\\\ 2.80 &   6870 $\\pm$ 263 &  24000 $\\pm$ 490\\\\ 3.00 &   2020 $\\pm$ 143 &  19100 $\\pm$ 441\\\\ \\hline \\hline \\hline \\end{tabular} \\end{center}  \\newpage \\bigskip \\centerline{} \\bigskip  \\centerline{\\bf Table 3} \\medskip \\centerline{\\bf Cross Sections for KBO Scattering} \\bigskip \\centerline{[all cross sections in units of (AU)$^2$]} \\bigskip  \\begin{center} \\begin{tabular}{cccccc} \\hline \\hline \\hline $\\epsilon$ & $a$ = 30 AU & $a$ = 40 AU & $a$ = 50 AU & $a$ = 60 AU & $a$ = 70 AU \\\\ \\hline \\hline 0.05 &172000 $\\pm$ 1670 &206000 $\\pm$ 1780 &236000 $\\pm$ 1860 &263000 $\\pm$ 1780 & 286000 $\\pm$ 1820\\\\ 0.10 &147000 $\\pm$ 1580 &179000 $\\pm$ 1700 &206000 $\\pm$ 1780 &229000 $\\pm$ 1720 & 251000 $\\pm$ 1770\\\\ 0.20 &119000 $\\pm$ 1450 &145000 $\\pm$ 1570 &169000 $\\pm$ 1670 &189000 $\\pm$ 1630 & 209000 $\\pm$ 1680\\\\ 0.30 & 99600 $\\pm$ 1350 &123000 $\\pm$ 1480 &143000 $\\pm$ 1570 &161000 $\\pm$ 1540 & 179000 $\\pm$ 1600\\\\ 0.40 & 86700 $\\pm$ 1270 &106000 $\\pm$ 1390 &126000 $\\pm$ 1500 &140000 $\\pm$ 1460 & 156000 $\\pm$ 1530\\\\ 0.50 & 76100 $\\pm$ 1200 & 92800 $\\pm$ 1310 &110000 $\\pm$ 1420 &123000 $\\pm$ 1390 & 138000 $\\pm$ 1460\\\\ 0.60 & 66700 $\\pm$ 1120 & 82900 $\\pm$ 1250 & 98600 $\\pm$ 1350 &109000 $\\pm$ 1320 & 123000 $\\pm$ 1400\\\\ 0.70 & 59600 $\\pm$ 1070 & 74100 $\\pm$ 1190 & 88300 $\\pm$ 1300 & 97800 $\\pm$ 1270 & 111000 $\\pm$ 1350\\\\ 0.80 & 52000 $\\pm$ 1000 & 66400 $\\pm$ 1130 & 78800 $\\pm$ 1240 & 87300 $\\pm$ 1210 &  99400 $\\pm$ 1290\\\\ 0.90 & 46000 $\\pm$  947 & 59100 $\\pm$ 1080 & 70400 $\\pm$ 1180 & 77800 $\\pm$ 1160 &  88900 $\\pm$ 1230\\\\ 1.00 & 39800 $\\pm$  889 & 52700 $\\pm$ 1030 & 61800 $\\pm$ 1110 & 67700 $\\pm$ 1090 &  78000 $\\pm$ 1170\\\\ \\hline \\hline \\hline \\end{tabular} \\end{center}   \\newpage"
    },
    "0011/astro-ph0011110_arXiv.txt": {
        "abstract": "We report the discovery of a new class of low frequency quasi--periodic variations of the X--ray flux in the X-ray bursters 4U1608-52 and 4U1636-536. We also report an occasional detection of a similar QPO in Aql X-1. The QPOs, associated with flux variations at the level of percents, are observed at a frequency of 7--9 $\\cdot 10^{-3}$ Hz. While usually the relative amplitude of flux variations increases with energy, the newly discovered QPOs are limited to the softest energies (1--5 keV). The observations of 4U1608-52 suggest that these QPOs are present only when the source X-ray luminosity  is within a rather narrow range and they disappear after X-ray bursts. Approximately at the same level of the source luminosity, type I X--ray bursts cease to exist.  Judging from this complex of properties, we speculate that a special mode of nuclear burning at the neutron star surface is responsible for the observed flux variations. Alternatively, some instabilities in the accretion disk may be responsible for these QPOs. ",
        "introduction": "Since the discovery of the quasi-periodic oscillations (QPO) in the X--ray flux from GX 5-1 (\\cite{vdk85}, Lewin, van Paradijs \\& van der Klis 1988), they were considered an important probe of the inner part of an accretion disk and a region where the accretion disk is interacting with a neutron star surface or magnetosphere (e.g \\cite{alpar85}, for a review of the present status of QPO observations and theoretical models see e.g. \\cite{vdk00}). The characteristic timescales in these regions are short and the QPOs are usually observed (and expected) at frequencies of tens or even thousands of Hz. Aperiodic variability is also present at much lower frequencies (e.g. $10^{-2}$--$10^{-3}$ Hz). In particular, when the spectrum of the accreting neutron star binaries is soft above $\\sim$ 5 keV (e.g. during the so-called atoll banana state) a Very Low Frequency Noise (VLFN) is observed at these frequencies (e.g. \\cite{hasingervdk89}), with an approximately power law dependence of power on frequency and typical RMS variations at the level of percents. Using RXTE data on several accreting neutron star binaries, we searched for low frequency quasi--periodic variations of the X--ray flux and found clear signatures of QPOs with a surprisingly similar frequency in 3 sources. The properties of the newly-found low frequency QPOs are very distinct from those of ``canonical'' high frequency QPOs, probably indicating a different underlying physics.  We describe the experimental results in Section 2. In Section 3 we speculate on a possible origin of the low frequency QPOs. Section 4 summarizes our findings. ",
        "conclusions": "A new type of very low frequency QPO  has been discovered in at least  two ``atoll'' sources: 4U1608-52 and 4U1636-536, with essentially the same frequency $f\\sim$ 7--9 $\\cdot 10^{-3}$ Hz. These QPOs have a very soft spectrum, RMS amplitude of variations of the order of per cent and appears to be present only at a certain level of the mass accretion rate. We suggest that these quasi--periodic variations might be related to a special regime of the nuclear burning on the neutron star surface or some instability in the optically thick accretion disk."
    },
    "0011/astro-ph0011056_arXiv.txt": {
        "abstract": "We use double-mode RR Lyrae (RRd) stars from the {\\sc macho} variable star database of the Large Magellanic Cloud (LMC) to estimate its distance, by utilizing photometric data, linear pulsation and stellar atmosphere models. If we set $E_{B-V}=0.11$ and [M/H]$=-1.5$ for LMC, we get $M-m=18.52$~mag, which is in agreement with the value obtained from Galactic cluster RRd stars. Although the formal statistical error (standard deviation) of $M-m$ derived from the {\\sc macho} set is only $0.024$~mag, the systematic errors are more substantial. These errors are mostly due to the zero point ambiguity of the temperature scale and to the error of the photometric calibration. However, the fact that the distance moduli obtained in this and in our former studies of cluster RRd and Small Magellanic Cloud beat Cepheids agree so well, implies that the only serious source of error is the zero point of the temperature scale, which should not have larger than $\\pm 0.10$~mag effect on the distance modulus. ",
        "introduction": "The determination of the distance moduli ($DM$) of the Magellanic Clouds plays an important role in establishing the cosmic distance scale. Mostly due to the analyses of the {\\sc hipparcos} data, there is a renewed effort to pin down this crucial distance within an accuracy of better than $\\approx 0.05$~mag. However, the question is far from being settled, and the LMC distance moduli obtained with various methods may differ by several tenth of magnitudes, stretching from the `short' ($DM\\approx 18.2$, e.g., Stanek et al. 2000) to the `long' ($DM\\approx18.6$, e.g., Feast 1999) distance scales.  In the first paper of this series (Kov\\'acs \\& Walker 1999, hereafter KW99, see also Kov\\'acs 2000b) we used RRd stars in Galactic clusters to show that they yield $0.2-0.3$~mag larger distances than the ones inferred from the Baade-Wesselink absolute magnitudes of RR~Lyrae stars. This result leads to $DM\\approx 18.5$~mag for the LMC. When the same method is applied to the Small Magellanic Cloud (SMC) beat Cepheids with 0.51~mag relative distance between the SMC and LMC, we get almost the same $DM$ for the LMC (Kov\\'acs 2000a). In the present Letter we directly extend the analysis to the {\\sc macho} RRd stars (Alcock et al. 1997, 2000) and show that these give larger distance as well, in fine agreement with the above mentioned independent data sets. ",
        "conclusions": "The recent revision of the {\\sc macho} RRd magnitudes (Alcock et al. 1999) enabled us to calculate the distance of the LMC directly from the observed magnitudes, colors and periods. This shows a very good agreement with the independent estimate based on Galactic globular cluster RRd stars and SMC beat Cepheids (Kov\\'acs 2000a). They all give $M-m\\approx18.5$~mag, which lends further support to the `long' distance scale. The largest source of systematic error is the potential ambiguity in the zero point of the temperature scale. Even this is not likely to cause the large effect, which would be necessary to lower the distance modulus by $\\approx 0.3$~mag, the value required by some current estimates based on other methods (e.g., Stanek et al. 2000). This amount of change would demand about $300$K lower temperature scale than the one used in this work. Since our adopted scale is only 100~K hotter than the presently published lowest infrared flux scale of Alonso et al. (1999), we think that it is highly unlikely that the required lowering of the distance modulus is possible by temperature change alone. Consequently, we would need some `conspiracy' with metallicity, reddening or other parameters, so as to get the necessary decrease in the distance modulus. This situation is also not very probable, because, although the strongest effect comes from the metallicity, this is primarily differential i.e., the distance moduli derived from the low metallicity clusters would not be seriously affected. This would lead to inconsistency among the distance moduli, because the relative distances, being empirically determined, are independent from the adopted metallicities. It is also worthwhile to mention that the metallicity values used in this and in our previous studies are in very good agreement with the ones published by independent spectroscopic studies and that these metallicities minimize the differences between the LMC distance moduli derived from various double-mode stars.  We conclude that the application of standard temperature and metallicity scales leads to the `long' RR~Lyrae distance scale, which, except for a drastic lowering of the temperature scale, is not possible to match with the `short' distance scale."
    },
    "0011/astro-ph0011042_arXiv.txt": {
        "abstract": "I give a brief review of high energy gamma-ray signatures of dark matter.  The decay of massive $X$-particles and subsequent hadronization have been suggested as the origin of the highest energy cosmic rays.  Propagation over cosmological distances to Earth (as would be the case in some topological defect origin models for the $X$-particles) results in potentially observable gamma-ray fluxes at GeV energies.  Massive relic particles on the other hand, would cluster in galaxy halos, including that of our Galaxy, and may give rise to anisotropic gamma ray and cosmic ray signals at ultra high energies.  Future observations above 100 Gev of gamma rays due to WIMP annihilation in the halo of the Galaxy may be used to place constraints on supersymmetry parameter space. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011568_arXiv.txt": {
        "abstract": "In \\Nbd\\ simulations of collisionless stellar systems, the forces are softened to reduce the large fluctuations due to shot noise. Softening essentially modifies the law of gravity at $r=|\\bx_i-\\bx_j|$ smaller than some softening length $\\epsilon$. Therefore, the softened forces are increasingly biased for ever larger $\\epsilon$, and there is some optimal $\\epsilon$ which yields the best compromise between reducing the fluctuations and introducing a bias. Here, analytical relations are derived for the amplitudes of the bias and the fluctuations in the limit of $\\epsilon$ being much smaller than any physical scale of the underlying stellar system. In particular, it is shown that the fluctuations of the force are generated locally, in contrast to the variations of the potential, which originate from noise in the whole system. Based on these asymptotic relations and using numerical simulations, I study the dependence of the resulting force error on the number of bodies, the softening length, and on the functional form by which Newtonian gravity is replaced. The widely used Plummer softening, where each body is replaced by a Plummer sphere of scale radius $\\epsilon$, yields significantly larger force errors than do methods in which the bodies are replaced by density kernels of finite extent. I also give special kernels, which reduce the errors even further. These kernels largely compensate the errors made with too small inter-particle forces at $r<\\epsilon$ by exceeding Newtonian forces at $r\\sim\\epsilon$. Additionally, the possibilities of locally adapting $\\epsilon$ and of using unequal weights for the bodies are investigated. By using these various techniques without increasing $N$, the rms force error can be reduced by a factor $\\sim2$ when compared to the standard Plummer softening with constant $\\epsilon$. The results of this study are directly relevant to \\Nbd\\ simulations using direct summation techniques or the tree code for force evaluation. ",
        "introduction": "\\label{sec:intro} \\Nbd\\ techniques are applied to stellar dynamical problems ranging from collision-dominated systems, like stellar clusters and galactic centres, to systems with extremely large numbers of particles, such as galaxies and large-scale structure. In the first case, one is actually interested in simulating a system of $N$ gravitationally interacting particles, which is done by {\\em collisional\\/} \\Nbd\\ codes.  In systems with very large numbers of particles, on the other hand, the two-particle relaxation time greatly exceeds the age, i.e.\\ these systems behave essentially {\\em collisionless,} very much like a continuous system. In other words, the shot noise is negligible, and each particle essentially moves in the mean force field generated by all other particles.  Thus, the aim of any collisionless \\Nbd\\ simulation is to simultaneously solve the collisionless Boltzmann equation (CBE) \\begin{equation} \\label{cbe} {\\partial f\\over\\partial t} + {\\partial f\\over\\partial\\bx}\\cdot\\bv + {\\partial f\\over\\partial\\bv}\\cdot\\bF = 0, \\end{equation} describing the stellar dynamics, and \\begin{equation} \\label{grav} \\bF(\\bx,t) = -G \\int\\dx^\\prime\\;\\int\\dv\\; {\\bx-\\bx^\\prime\\over|\\bx-\\bx^\\prime|^3}\\,f(\\bx^\\prime,\\bv,t) \\end{equation} describing gravity. Here, $f(\\bx,\\bv,t)$ is the (continuous) one-particle% \\footnote{Correlations between particles (e.g.\\ binaries) are not described by $f$ but by the two-particle distribution function and higher orders in the BBGKY hierarchy (see, e.g., Binney \\& Tremaine 1987).} phase-space distribution function (DF) of the stellar system, while $\\bF(\\bx,t)$ is the gravitational force field generated by it. A numerical treatment of the CBE on an Eulerian grid is not feasible because of the vast size of six-dimensional phase space and the strong inhomogeneity of the DF. Instead, one applies the method of characteristics. Since the CBE states that $f$ is constant along any trajectory $\\{\\bx(t),\\bv(t)\\}$, the trajectories obtained by time-integration of $N$ points $\\{\\bx_i,\\bv_i\\}$ sampled from the DF at $t=0$ form a representative sample of $f$ at each time $t$.  The result of this procedure is a \\Nbd\\ code, however, with a crucial difference to collisional \\Nbd\\ codes: the bodies do not represent real particles, rather they are a Monte-Carlo representation of the underlying continuous DF. In fact, all the information on $f(\\bx,\\bv,t)$ is given by the phase-space positions of the bodies. Since $N$ is much smaller than the number of particles in the system being modelled, this information is always incomplete.  \\subsection{The rationale of force softening} \\label{sec:intro:why} Having solved for the dynamics, we still need to compute gravity. The simple-minded approach of replacing $f$ in \\eqn{grav} with a sum of $\\delta$-peaks at the body positions corresponds to a Monte-Carlo integration \\cite{pp93} and yields the forces of $N$ mutually interacting particles.  In practice, one {\\em softens\\/} the forces at small inter-body separations $r=|\\bx^\\prime-\\bx|$, which is equivalent to replacing the bodies by some extended mass distribution instead of $\\delta$-functions\\footnote{ In describing their dynamics, on the other hand, the bodies are treated as {\\em point-like\\/} test particles. Dyer \\& Ip (1993) argued that this asymmetry is physically inconsistent. However, bodies do not represent physical particles and no physical inconsistency can arise. The goal is {\\em not\\/} the simulation of $N$ mutually interacting particles point-like or not.}. Since there appears to be some confusion for the proper reasons to introduce softening, I will now try to give a thorough motivation.  \\subsubsection{Suppressing artificial relaxation? No!} \\label{sec:intro:relax} Contrary to widespread belief, softening does not much reduce the (artificial) two-body relaxation \\cite{theis98}, the time-scale of which scales like $N/\\ln N$ and which cannot be neglected in {\\em simulations\\/} of collisionless stellar systems.  To see this, recall the works of Chandrasekhar and Spitzer, who showed that two-body relaxation is driven by close as well as distant encounters: each octave in distance is contributing equally. By softening the forces at small $r$, one reduces the contributions from close encounters. Most of the relaxation, however, is due to noise on larger scales, and hence cannot be avoided by softening techniques, in agreement with the numerical findings of Hernquist \\& Barnes (1990). \\subsubsection{Avoiding artificial two-body encounters} \\label{sec:intro:avoid} Without softening, the accurate numerical integration of close encounters and binaries, the main obstacle in collisional \\Nbd\\ codes, requires great care and substantial amounts of computer time. However, as already mentioned, the bodies are just a representation of the {\\em one-particle\\/} DF, and consequently binaries as well as two-body encounters (and the Newtonian forces arising in them) are entirely artificial. Thus, softening may be motivated by the desire to suppress such artifacts. \\subsubsection{Estimating the true forces} \\label{sec:intro:estim} One may consider force softening as a way to obtain, at each time step, from the $N$ body positions an {\\em estimate\\/} for the gravitational forces generated by the continuous mass density of the system being modelled (Merritt 1996, see also below). In this context, softening essentially implies the assumption that the mass density is smooth on small scales. This assumption compensates for the aforementioned incompleteness in our knowledge of the DF, a standard method in non-parametric estimation techniques. \\subsubsection{The benefits of softening} As a benefit of softening, close encounters between bodies are much less of a problem: one can use the simple leap-frog integrator, a method that would be highly inefficient in collisional \\Nbd\\ codes. Moreover, since only an estimated force fields is computed, approximate, and hence cheaper, methods for its computation than direct summation may be used, as long as the errors introduced by the approximation are smaller than those due to the estimation itself. Both these effects lead to substantial savings of computer time, which in turn enable much larger numbers of bodies in collisionless than in collisional \\Nbd\\ codes (currently about 3-4 orders of magnitude more). This may be considered the real motivation for softening.  \\subsection{Optimal softening} Softening also has a severe drawback: it artificially modifies mutual gravitational forces at small inter-body separations: instead of Newton's inverse square law, the forces vanish in the limit $r\\to0$ (necessary to yield a continuous $\\bF(\\bx)$). With increasing degree of softening, this results in an ever stronger {\\em bias\\/}, the systematic error of the mean estimated as compared to the true forces, corresponding to a diminution of the information hold in the body positions.  Thus, softening allows a trade-off between artificially high near-neighbour forces and a biased force field. Both do affect the fidelity but also the efficiency of a \\Nbd\\ simulation. With little softening, the integration of the equations of motion requires great care, demanding substantial computer resources (if the time integration is not done carefully enough, two-body relaxation may be strongly enhanced and other unwanted effects may arise). On the other hand, the force bias introduced by softening modifies the dynamics, possibly on scales much larger than that on which forces are softened, which might deteriorate the simulation. The influence of these two effects on a \\Nbd\\ simulation will certainly depend on (i) the purpose of the simulation, (ii) the properties of the system modelled and (iii) the time span integrated. This implies that the optimal way of softening depends on the nature of the problem investigated as well as the goal of the simulation.  The idea of this paper is to make a first step towards answering the question of the optimal softening by investigating the effect of softening on the estimated force. That is, here we consider only the {\\em static\\/} problem of optimal force estimation. An investigation of the effect of softening on the {\\em dynamics\\/} is beyond the scope of this paper, but see the discussion in Section~\\ref{sec:disc:eopt}.  In some situations, one can reduce the noise by a careful rather than random choice of the trajectories, the so-called ``quiet start'' technique \\cite{se83}. However, if chaotic motion dominates most of the trajectories, e.g.\\ in a merger simulation, the benefits of such a careful set-up quickly disappear. Here, I am interested in the general, i.e.\\ worst, case, and assume throughout that the trajectories are sampled randomly from the DF.  \\subsection{A formalism of softening} A softening technique that is widely used is the Plummer softening, initiated by Aarseth's (1963) use of Plummer spheres to model galaxies in a cluster. In this method, each body is replaced by a Plummer sphere with scale radius $\\epsilon$, resulting in the potential estimator (hereafter, an estimate for some quantity is indicated by a hat) \\begin{equation} \\label{plummer-soft} \\hPhi(\\bx) =  -G \\sum_{i=1}^N {m\\over\\sqrt{\\epsilon^2+(\\bx-\\bx_i)^2}}, \\end{equation} where $m=M/N$ is the weight (or mass) of each individual body. A more general estimator for the gravitational potential is \\begin{subequations} \\label{kern} \\begin{equation} \\label{soft:pot} \\hPhi(\\bx) = -G \\sum_{i=1}^N {m\\over\\epsilon}\\; \\phi\\left[{|\\bx-\\bx_i|\\over\\epsilon}\\right]. \\end{equation} This formula clearly separates the two aspects of the softening method: the {\\em softening kernel\\/} $\\phi(r)$, which determines the functional form of the modified gravity, and the {\\em softening length\\/} $\\epsilon$. The Plummer softening, for instance, corresponds to $\\phi=(1+r^2)^{-1/2}$. The corresponding estimates for the force field $\\bF=-\\bnabla\\Phi$ and the density $\\rho$ of the underlying smooth system are \\begin{eqnarray} \\label{soft:force} \\hbF(\\bx) &=&-G \\sum_{i=1}^N {m\\over\\epsilon^2}\\, \\phi^\\prime\\left[{|\\bx-\\bx_i|\\over\\epsilon}\\right]\\, {\\bx-\\bx_i\\over|\\bx-\\bx_i|},                    \\\\ \\label{soft:density} \\hrho(\\bx) &=& \\sum_{i=1}^N {m\\over\\epsilon^3}\\, \\eta\\left[{|\\bx-\\bx_i|\\over\\epsilon}\\right], \\end{eqnarray} \\end{subequations} where \\begin{equation} \\eta(r) =  -{1\\over4\\pi r^2}\\, {\\partial\\over\\partial r} \\Big(r^2\\, {\\partial\\phi\\over\\partial r}\\Big), \\end{equation} is the {\\em kernel density}, the mass distribution by which each body is replaced in these estimates. Note that kernels $\\phi(r)$ which reproduce Newtonian gravity exactly for $r$ larger than some finite radius $r_0$ correspond to density kernels of finite support, i.e.\\ $\\eta(r)\\equiv0$ for $r>r_0$.  Following the terminology of statisticians (e.g..\\ Silverman 1986), the above approximations in \\eqns{kern} may be called {\\em fixed-kernel estimators\\/} because the softening length $\\epsilon$ is the same for all bodies.  \\subsection{Outline of the paper} The reader might be surprised by this introduction being rather void of quotations. In fact, even though \\Nbd\\ methods with softening are widely used by astronomers, no detailled investigation of the properties of force softening has yet been undertaken. Some papers deal with the small-scale truncation of the mutual forces, but only few consider this bias in conjunction with the reduction of noise. Merritt \\& Tremblay (1994) dealt with estimators for the density of stellar systems, and pointed out the close connection to softening in \\Nbd\\ simulations. Merritt (1996) introduced the important concept of the mean square force error (see Section~\\ref{sec:errors} below) and studied empirically the effect of the softening length on the errors of the forces estimated by Plummer softening.  Athanassoula et~al.\\ (1998, 2000) have extended this work to larger $N$ and to using the tree code rather than direct summation as Poisson solver.  A systematic investigation of the dependence of the force errors on the softening method, in particular for other than Plummer softening, is still missing, and the goal of this paper is to fill this gap. Merritt and Athanassoula et~al.\\ found, by brute force, empirical relations for the dependence on $N$ of the optimal softening length, $\\epsilon_{\\rm opt}$, and the resulting force error, mainly for the case of Plummer softening. Instead, I derive in Section~\\ref{sec:errors} and Appendix~\\ref{app:asym} of this paper analytic expressions, valid for softening of the form \\eqb{kern}, for the force error in the asymptotic limit of small $\\epsilon$ and large $N$. These asymptotic relations apply whenever $\\epsilon$ is smaller than any physical scale of the stellar system being modelled, a condition that must be satisfied for the \\Nbd\\ simulation to be of any use. Based on the results of this investigation, I discuss in Section~\\ref{sec:opt} the optimal softening length and kernel, and the asymptotic scaling with $N$ of $\\epsilon_{\\rm opt}$ and the error achieved, while numerical experiments that demonstrate and verify these analytic results are presented in Section~\\ref{sec:exper}.  Apart from just increasing the number $N$ of bodies and of the choice of the softening kernel and length, the \\Nbd\\ experimenter may use a scheme to locally adapt the softening, and use different individual weights. These two latter techniques, which aim at enhancing the resolution in regions of higher density and thus decreasing the force error, are discussed in Section~\\ref{sec:ind} and related numerical experiments are presented in Section~\\ref{sec:exper:adaptive}. In Section~\\ref{sec:disc}, I discuss the relevance of the results of these investigations to other Poisson solvers than direct summation, and touch the question for the optimal collisionless \\Nbd\\ code and related issues. Finally, Section~\\ref{sec:summ} sums up and concludes. ",
        "conclusions": "\\label{sec:summ} In collisionless \\Nbd\\ simulations, the bodies are the numerical representation of the continuous one-particle DF. Because $N$ is much smaller than in the stellar system being modelled, this representation is always incomplete, i.e.\\ noisy. In order to estimate, at each time step, the continuous force field due to the DF, one has to moderate this noise. This exactly is the purpose of force softening in collisionless \\Nbd\\ codes. However, while reducing the noise in the \\Nbd\\ forces, the softening modifies the laws of gravity: the force vanishes at vanishing inter-body separation $r\\to0$. This generates a bias, a systematic offset between the estimated and true forces. Both, the noise and the bias, contribute to the mean square error of the estimated force field. While the noise decreases with the softening length $\\epsilon$, the bias increases, and the optimal softening represents a compromise between these two \\cite{merritt96}.  In the limit of $\\epsilon$ being small compared to any physical scale of the stellar system, a requirement that is desirably satisfied in any \\Nbd\\ simulation, analytic expressions for the noise-generated variance of the force and for the bias have been derived. While the relations \\eqb{bias} for the bias have already been given in the literature (e.g.\\ Hernquist \\& Katz 1989), those for the variance of the estimated potential and force \\eqi{vars} seem to be hitherto unknown. One of the most important results is that the variance of the force is a local quantity diverging in the limit $\\epsilon\\to0$, very much like the variance of the density, while the variance of the potential depends on a global measure and reaches a finite value for $\\epsilon=0$. Presumably, this is related to the fact, see Section~\\ref{sec:disc:relax} and Theis (1998), that the two-body relaxation is not substantially reduced by force softening.  In these asymptotic relations, the effects of the softening length, the softening kernel, the number $N$ of bodies, and of the underlying stellar system nicely separate, allowing simple and, in the important case of large $N$, accurate estimates for the behaviour of the mean integrated squared force error, \\miseF.  For the various types of softening kernels, Table~\\ref{tab:asym} gives the asymptotic relations for the optimal softening length and force errors (the relations given for Plummer softening are based on numerical experiments only). The strongest possible dependence of the MISE on $N$ is $N^{-1}$ as pertains, e.g., to parametric estimators.  Note that the compensating kernels come very close to this limit. The main results of this investigation are: \\setcounter{lll}{0} \\begin{list}{\\arabic{lll}.}{\\usecounter{lll} \\leftmargin4mm \\itemsep0pt plus 1pt \\parsep0pt plus 1pt \\topsep2pt plus 1pt \\labelwidth3mm \\labelsep1mm \\itemsep2ex} \\item   Compact softening kernels, i.e.\\ those with finite density support, are superior to the standard Plummer softening. This improvement is ever larger for increasing $N$. \\item   Special softening kernels are derived, for which the force errors are even smaller. These kernels compensate the small forces at $r\\to0$ by forces larger than Newtonian at $r\\sim\\epsilon$. \\item   For inhomogeneous stellar systems, such as galaxies, adaptive softening lengths allow a further substantial reduction of the force errors. \\end{list} All these numbers are based on the Hernquist model, which gives a fair representation of a typical galaxy, as underlying stellar system, resulting in a rms force error of $\\sim0.01=0.04$ of the rms force of this model. For smaller errors to be achieved, the advantage of using compensating kernels and adaptive softening lengths will become even larger.  For this error of 4\\%, particle numbers $N\\sim10^5$ are necessary, and roughly 20 times as many to reduce the error to 1\\%. It is not clear, what these errors imply in terms of the reliability of the time evolution of a corresponding \\Nbd\\ simulation. This is a very important question to answer, in order to better understand the results of \\Nbd\\ simulations and assess their reliability. For example, current \\Nbd\\ simulations of large-scale structure formation employ only about $10^{3-4}$ bodies per halo, which come out to have a profile similar to a Hernquist model, i.e.\\ the rms force error is $10\\%$ to $25\\%$, even if $\\epsilon$ was chosen optimally, which it very likely was not."
    },
    "0011/astro-ph0011274_arXiv.txt": {
        "abstract": " ",
        "introduction": "A4059 at $\\sim300$ Mpc (z = 0.0487) is well studied in the X-ray, radio, and optical.  It is relatively compact and is classified as richness class l (Schwartz et al., 1991 and references therein).  It is dominated by a central CD galaxy, ESO 349-G010, and contains the radio galaxy PKS 2354-35 (Taylor, Barton, \\& Gee, l994). Taylor et al. (1994) found that the radio galaxy has a double lobe structure extending along the major axis of the host galaxy. X-ray observations with HEAO I and EXOSAT (Edge \\& Stewart 1991; Schwartz et al. 1991; Edge, Stewart \\& Fabian 1992) showed the cluster has substantial X-ray cluster emission with a strong cooling flow.  ASCA observations (Ohashi 1995) found a chemical abundance gradient in the cluster gas around the central CD galaxy.  Huang \\& Sarazin (1998) analyzed ROSAT, HRI and PSPC X- ray data on the cluster and found two X-ray holes near the center of the cluster emission.  We have studied A4059 in the hopes that it might shed light on the underlying source mechanism for EUV emission in clusters of galaxies.  It is a particularly interesting cluster to study in this regard because it appears to have achieved a relaxed state (Huang \\& Sarazin 1998) and has a cooling flow.  These characteristics are similar to those in A1795 and A2199 and Bowyer et al. (1999) showed these clusters did not exhibit excess EUV emission. On the other hand, A4059 does contain a radio galaxy, and Bergh\\\"ofer et al. (2000) have shown that radio emission may be at least indirectly associated with the EUV emission in the Virgo Cluster.  We report herein our analysis of the EUV emission from A4059. In view of recent claims that the EUVE Deep Survey Telescope detector background is uncertain since it varies over time, and that the images obtained are highly dependent upon the threshold setting employed (Lieu et al. 1999), we have examined these issues in detail. Throughout this paper we assume a Hubble constant of 50 Km\\,s$^{-1}$\\,Mpc$^{-1}$ and q(0) = 0.5. ",
        "conclusions": "We analyzed unpublished EUVE Deep Survey observations of the cluster of galaxies A4059. In order to test the integrity of our results for this clusters and to test the validity of our background model for the EUVE Deep Survey instrument we also analyzed 363\\,Ks of blank field observations and compared this to a previous assemblage of 425\\,Ks of blank sky. Using statistical tests of these two backgrounds we definitely confirm the stability of the instrument's background and disprove claims of time dependent changes in the detector background by Lieu et al. (1999). Further analysis shows that for a wide range of threshold cutoffs the pulse-height threshold level setting has essentially no effect on the final EUVE Deep Survey images, another claim cooked up by Lieu et al. (1999).  Our analysis of the EUVE observations of the cluster of galaxies A4059 shows that this cluster exhibits diffuse EUV emission. However, the emission over most of the cluster is that expected from the low energy contribution of the X-ray emitting cluster gas. The EUV emission in the central 2 arcmin of the cluster shows an EUV deficit. The observed level of EUV flux is extremely sensitive to absorption effects and the flux deficit demonstrates intrinsic absorption of the X-ray emitting cluster gas. Together with A1795 and A2199 (Bowyer et al. 1999), A4059 is the third example of a cluster of galaxies with EUV absorption due to cooler gas in the existing central cooling flow."
    },
    "0011/astro-ph0011104_arXiv.txt": {
        "abstract": "We describe a new approach to obtaining very high surface densities of optical spectra in astronomical observations with extremely accurate subtraction of night sky emission. The observing technique requires that the telescope is nodded rapidly between targets and adjacent sky positions; object and sky spectra are recorded on adjacent regions of a low-noise CCD through charge shuffling. This permits the use of extremely high densities of small slit apertures (`microslits') since an extended slit is not required for sky interpolation. The overall multi-object advantage of this technique is as large as 2.9$\\times$ that of conventional multi-slit observing for an instrument configuration which has an underfilled CCD detector and is always $>1.5$ for high target densities. The `nod-shuffle' technique has been practically implemented at the Anglo-Australian Telescope as the ``LDSS++ project'' and achieves sky-subtraction accuracies as good as 0.04\\%, with even better performance possible.  This is a factor of ten better than is routinely achieved with long-slits. LDSS++ has been used in various observational modes, which we describe, and for a wide variety of astronomical projects. The nod-shuffle approach should be of great benefit to most spectroscopic (\\eg long-slit, fiber, integral field) methods and would allow much deeper spectroscopy on very large telescopes (10m or greater) than is currently possible. Finally we discuss the prospects of using nod-shuffle to pursue extremely long spectroscopic exposures (many days) and of mimicking nod-shuffle observations with infrared arrays. \\begin{center} {\\it Accepted for publication in PASP} \\end{center} ",
        "introduction": "The problem of subtracting the night sky foreground emission is a critical one for astronomical spectroscopy. The task is particularly acute in the red part of the spectrum (600$-$1000nm) as there are numerous hydroxyl (OH) bands which dominate the light giving a bright background. Many authors have recognized over the past twenty years that low to moderate resolution spectroscopy in this band is ultimately limited by systematic uncertainty associated with sky subtraction (\\eg Dressler 1984).  In some respects, it is surprising that optical astronomy has been slow to recognize an important technique utilized by near-infrared astronomy, \\ie beam-switching. Here, the background signal is very strong, is highly variable, and influences all observations (\\eg Ramsay, \\etal, 1992). A common implementation of beam-switching is where the secondary mirror `chops' between a target object and a sky field while the infrared array is read out continually.\\footnote{`Chopping' refers to a moving secondary mirror while the primary remains fixed on the object; we use `nodding' to indicate a fixed secondary where the pointing of the primary mirror alternates between sky and an object field.}  This is perhaps because there is a conflict between the desire to beam-switch rapidly, and sample the sky contemporaenously, and the desire to take long integrations to minimize the effect of readout noise. This is especially true for modern, very low noise CCD detectors.  The underlying principle of the nod-shuffle technique is simply that a CCD detector can be used to store two images of a field, imaged quasi-simultaneously (Cuillandre\\etal\\ 1994; Bland-Hawthorn 1994; Sembach \\& Tonry 1996). By using `charge-shuffling' charge can be moved from an  illuminated region to a storage region. This process does not invoke readout noise and only takes only a fraction of a second since charge can be shifted between CCD rows two to three orders of magnitude faster than it can be read out. If this shuffling is synchronized to telescope motion two interleaved exposures of object and sky can be imaged side by side at the detector. Note three important facts: (i) the images are obtained through identical optical paths, (ii) the imposed flatfield structure is identical for both images, and (iii) the CCD is read out only once.  The use of shuffling techniques in astronomy can be traced to early attempts to improve the performance of imaging polarimeters (McLean\\etal\\ 1981; Stockman 1982).  Since that time, charge-shuffling has been little utilised.  Part of the reason may stem from experiments by Lemonier \\& Piaget (1983). By rapidly shifting charge backwards and forwards many times (pocket pumping), they were able to identify local defects in the potential profile (trapping sites) within the silicon substrate. By the end of the 1980s, traps and deferred charge were still a fundamental limitation to repeated charge shuffle operations (Blouke\\etal\\ 1988).  The development of charge-shuffling at the Anglo-Australian Observatory dates back to the 1994 Marseilles conference on imaging spectrographs (Comte \\& Marcelin 1995).  It was here that the first results of integral field spectrographs were presented, arguably the most important development in optical instrumentation in the past decade. It was clear, and remains true, that the fundamental limitation of this powerful technology is the difficulty of accurate sky subtraction (Bland-Hawthorn 1995).  Key developments in CCD manufacture have made charge-shuffling a realistic prospect and an important consideration in all future instrument design. First, the latest generation CCDs (EEV, MIT Lincoln Lab) have very low read noise ($\\sim 1e^-$), negligible dark current, high purity and very high charge transfer efficiency (99.9999\\%). Secondly, the manufacturing process prefers to generate rectangular arrays\\footnote{The photofab process uses a $25-30$mm reticle which restricts the `row' dimension of a CCD. The reticle is stepped down the wafer and the new circuit is stitched to the previous pattern.} which provide for storage regions. Bland-Hawthorn \\& Barton (1995) demonstrate that, with modern CCDs, it takes more than a hundred shuffle operations before bulk trapping sites start to compromise the data.  In this paper, we describe the development at the AAO of the `nod-shuffle' method founded on the principle of CCD charge shuffling.  This differential technique has resulted in two important experimental breakthroughs. First, the object and sky can be measured quasi-simultaneously. As we show, the main limit to the accuracy of sky-subtraction is the rapidity of nod-shuffling compared to the temporal power spectrum of sky brightness variations. Secondly, nod-shuffle allows for a considerable increase in the {\\it multi-object gain} of a spectrograph,  up to 2.9$\\times$ more objects per unit observing time using small `microslits,' for fields with high object densities. We have implemented the nod-shuffle method with the Low Dispersion Survey Spectrograph (LDSS) on the Anglo-Australian Telescope (AAT) and have obtained fractional residuals as low as $4\\e{-4}$.  The plan of this paper is as follows: in Section~2 we describe the nod-shuffle concept and discuss qualitatively the sky-subtraction and multiplex advantages to be gained. In Section~3 we describe in detail our implementation of nod-shuffle at the AAT using the Low Dispersion Survey Spectrograph and show some example data. In Section~4, we show the increased multi-object gain which becomes possible via the nod-shuffle operation. We quantify the sky-subtraction accuracy in Section~5 and discuss ways in which it might be improved further.  In Section~6, we illustrate key observing modes for LDSS++ which are facilitated by the use of microslits.  Finally, we discuss future prospects for the nod-shuffle observing mode. ",
        "conclusions": "We have explained the virtues of the nod-shuffle technique for CCD-based optical spectroscopy: we reach a new level of sky-subtraction precision of 0.04\\%. This is in accord which predictions from a reasonable physical model of atmospheric airglow.  This technique also permits a great increase in the multiplex gain of multi-slit spectrographs we have quantified those gains and showed that they are the most in high-object density regimes.   We have outlined our thoughts on IR techniques equivalent to nod-shuffle. Possibly new circuit designs would allow charge storage but would need to be developed. Given the importance of IR spectroscopy on future large telescopes the scientific case for doing so is strong. Failing this we have outlined a less satisfactory, but still useful concept, for physically moving the IR array.   For very large telescopes (10m and greater) the precision of sky-subtraction is a real barrier for ultra-deep spectroscopic exposures. The systematic limit of ordinary slit subtraction is reached in only a few hours. The nod-shuffle technique offers a remedy, and promises the possibility of extremely long exposures, it's ultimate performance remains to be explored."
    },
    "0011/astro-ph0011332_arXiv.txt": {
        "abstract": "I give a review of catalogues of galaxies, clusters of galaxies and superclusters -- sources of information to study the large--scale structure of the Universe.  Thereafter I shall discuss the power spectrum of density perturbations, and the correlation function -- principal description functions which characterize the large--scale structure. I shall pay special attention to the geometric interpretation of these functions, i.e.\\ to the way in which the various properties of the distribution of galaxies in systems and systems themselves are reflected in these functions.  Finally, a discuss cosmological parameters which characterize general properties of the Universe -- the Hubble constant, densities of various populations of the Universe, and parameters of the power spectrum of galaxies and matter. ",
        "introduction": "In this review I shall accept certain paradigms on the structure of the Universe.  Paradigms on the structure of the Universe have evolved considerably during the last hundred years.  This process is continuing until the present time, and changes occur quite often, so it is important that we clearly state the present paradigms.  I shall use the term ``Universe'' for the real world around us, and the term ``universe'' for a model of the Universe (say Friedmann--universe). Our accepted paradigms are:  \\begin{itemize} \\item{} The Universe evolves from an explosive event termed ``Big Bang'', through the inflation, the radiation--domination era to the matter--domination era; probably we live now in the next era where the dominating constituent of the Universe is dark energy.  \\item{} The principal force driving the cosmological evolution is gravity.  \\item{} Density perturbations grow from small random fluctuations generated during the inflation; perturbations have Gaussian distribution.  \\item{} The main constituents of the Universe are: baryonic matter (stars and planets, hot and cold gas, and primordial gas in voids), dark matter, either cold (CDM) or hot (HDM), and dark (vacuum) energy.  \\item{} The Universe is flat -- the total mean density of all its populations is equal to the critical density.  \\end{itemize}  There exist a number of excellent reviews on the subject ``Large--scale structure of the Universe'', the most recent one with references to earlier work is the talk by Guzzo \\cite{g99} in the 19$^{th}$ Texas Symposium.  The term ``Large--scale structure of the Universe'' itself originated in the contemporary meaning as the title of an IAU Symposium \\cite{le78}, where the presence of filamentary distribution of galaxies and clusters with large empty voids between them was first reported.  In this review I use the experience collected at Tartu Observatory during the last 25 -- 30 years of the study of the Universe.  I shall pay attention to cosmographic aspects of the problem, in particular to the geometrical and physical interpretation of descriptive functions.  Also I try to show how advances in observational cosmology have changed our theoretical understanding of the formation and evolution of the structure of the Universe. ",
        "conclusions": ""
    },
    "0011/astro-ph0011454_arXiv.txt": {
        "abstract": "The isotropy in the ultra high energy cosmic ray (UHECR) flux observed by Yakutsk and AGASA experiments, is a very strong constraint to production and propagation models alike.  Most of the scenarios proposed in the literature should produce a sizable anisotropy as either extragalactic luminous or dark matter is normally associated with the invoked particle sources.  We explore the possibility that the magnetic fields in fossil cocoons of former radio galaxies -- so called {\\it radio ghosts} \\cite{Ringberg99}-- are able to scatter UHECR in the intergalactic medium giving rise to the observed isotropy. We show, through numerical simulations, under which conditions this process can be operative and the magnitude of the effect. We further demonstrate, that if radio ghosts mix with the ambient medium, they might be able to produce the observed magnetic fields in clusters of galaxies. In the case of mixing, the UHECR isotropization would be even stronger than in our conservative estimates. ",
        "introduction": "\\subsection{The Problem of UHECR Isotropy}  The upper end of the cosmic ray spectrum, at total energies above the (as yet unobserved) Greisen-Zatsepin-Kutzmin (GZK) cut-off ($\\sim 4 \\times 10^{19}$ eV - \\cite{GZK_G66}, \\cite{GZK_ZK66}) represents a challenge for particle-physicists and astrophysicists alike. The nature of the sources of these ultra-high energy cosmic rays and their distance scale are still unknown. Only our own galactic disk can be ruled out at present as a major source site, as a compatible anisotropy has not been observed by any of the experiments sensitive to the UHECR energy range \\cite{AG_1e19,Bird98,GMT_AAW_DM_halo}.  Although photons, neutrinos, or some unknown particle cannot be disregarded, the muon to electron ratio measured for extensive air showers \\cite{Halzen95,AGASA_pairs} points to hadrons as the primaries hitting the upper atmosphere and triggering the cascades. Neutrons with relativistic factor $\\gamma \\sim 10^{11}$ decay into protons after a path of $\\sim 1$ Mpc. Heavy nuclei, on the other hand, may lose as much as $\\sim 2-4$ nucleons per Mpc \\cite{Puget76,Cronin92} due to photodisintegration in interactions with the cosmic microwave and infrared backgrounds. Therefore, although new estimates of the infrared background \\cite{Malkan_and_Stecker_98} seem to indicate that the previous photodisintegration rates may be overestimated by a factor of $10$ \\cite{Stecker98}, UHECR, unless galactic, are very likely light nuclei, probably mainly protons. We will assume the latter in the remaining of this work.  Depending on the large scale configuration of the intergalactic magnetic field (IGMF), field values as low as $B_{IGMF} \\sim 10^{-9}$ G \\cite{Kromberg94}  may be expected. Consequently, proton gyroradius at $E\\sim 10^{20}$ eV can be of the order of $10^{2}$ Mpc. This represents a potentially interesting perspective for astronomy since UHECR should, under favorable conditions, point to their sources opening a new astronomical window. This would allow not only the direct study of their sources, but also the measurement of the intervening intergalactic and galactic (particularly halo) magnetic fields.  Relevant to this potential is the nature of the production mechanism involved. So far, both bottom-up and top-down UHECR production mechanisms have been devised and constitute perfectly viable scenarios. The first set of models invoke some kind of acceleration agent that is always spatially associated with baryonic matter, closely traced by luminous matter. Top-down mechanisms, on the other hand, resort to the decay of relics from GUT phase transitions in the early Universe \\cite{Sigl_TD}. Under the assumption that relics aggregate like dark matter (DM) and that their decay products are susceptible of inelastic interactions with the CMB, then any observed anisotropy should trace the nearby distribution of DM. Therefore, the UHECR flux should have two components, one originated inside the galactic halo and another extragalactic still roughly traced by luminous matter. \\cite{Dubovsky_DM} showed that, under general conditions, the halo component would dominate the extragalactic flux by at least two orders of magnitude. This is only true, however, in the unrealistic case of dark matter uniformly distributed in intergalactic space. Nevertheless, dark matter aggregates strongly and tends to be overabundant, by factors of $\\sim 10^{2}$, in the center of galaxy clusters when compared to its abundance in the halos of isolated galaxies. It can therefore be shown that \\cite{UHECRiso00}, in a sample of 47 events (as AGASA's above $4 \\times 10^{19}$ eV), and assuming Virgo as the only source of extragalactic events, 3-7 events should originate in Virgo and arrive inside a solid angle of approximately the size of the cluster. This could give rise to a slight anisotropy that correlates with the SGP when combined with the almost isotropic flux originated in a large galactic halo.  Photo-pion production interactions with the cosmic microwave background (CMB) restrict the source region to nominally $D \\stackrel{~}{<} 100$ Mpc \\cite{Waxman97,Hillas98}. Therefore, the signatures of the characteristic inhomogeneity of the local universe should be clearly observable in the angular arrival distribution of UHECR regardless of the production mechanism. In spite of the low statistics available at present, the latter does not seem to be the case \\cite{AG_1e19,GMT_isotropy00}.  If UHECR are charged particles originated inside the large scale structures in the nearby Universe, then any isotropization must be due to intervening magnetic fields. This could happen either inside the galactic halo, like in the presence of a magnetized galactic wind \\cite{Ahn_etal_GW00}, by the interaction with highly structured IGMF inside walls and filaments \\cite{GMT_APJL98_IGMF} or by the scattering off magnetic irregularities permeating the intergalactic medium.  In the present work we analyze the latter possibility considering radio ghosts \\cite{Ringberg99} as scattering centers of UHECR protons.      \\subsection{The Existence of Radio Ghosts} Active galaxies eject large amounts of radio emitting plasma -- short: {\\it radio plasma} -- into their environment. There it forms the typical cocoons of radio galaxies. The radio emission results from synchrotron emission of a population of relativistic electrons in the radio plasma's magnetic fields. It is possible to estimate the minimal energy density of the electron and magnetic components of the radio plasma required in oder to produce the observed emissivity.  This minimum is given by rough energy equipartition between electron population and the magnetic fields. The resulting minimal pressure is typically of the order of the environmental thermal pressure, indicating that relatively strong magnetic fields are present.  After a cosmological short time of $10^7 - 10^8$ years the radio luminosity of the cocoon decays strongly due to radiation and expansion energy losses of the electron population, or since the central engine of the radio galaxy stopped its activity.  Although undetectable by our instruments, the radio plasma is still present in the IGM inside a fossil radio cocoon, a so called radio ghost. The subsequent evolution of the radio plasma is unclear, since observationally poorly constrained. It can be expected that the strong magnetic fields of the ghosts allow it to resist erosion by subsonic turbulence. The fossil radio cocoon should be kinematically decoupled from the ballistic motion of the parent galaxy and follow mostly the flow pattern of the embedding material. From this one would expect the ghosts to have a cosmological distribution comparable to that of the galaxies. And one would further expect that the oldest ghosts are swept into clusters of galaxies by the flow of structure formation. But buoyancy of the probably very light radio plasma can produce some relative motion between ghosts and the IGM gas. That buoyancy is operating on radio plasma is impressively demonstrated by a recent simultion of a buoyant rising radio blob in a cluster center, which reproduces radio and X-ray observations of M87 strikingly well \\cite{churazov200sub}.   This might allow the ghosts in clusters or filaments of galaxies to ascend to larger radii, until they get stopped by freely infalling matter at the accretion shock.  It depends crucially on the topology of this accretion shock surface if the ghosts are able to escape from gravitationally bound structures as clusters and filaments of galaxies, or not. It is therefore difficult to predict the spatial distribution of ghosts. In this work we consider cases where the ghosts are distributed as the galaxies, where they are more, and where they are less clustered.  There are several observations supporting their existence and future experiments might be able to see further signatures of ghosts. \\begin{itemize} \\item The progenitor of radio ghosts, radio cocoons of active radio galaxies, are observed and well studied. Since the sharp boundaries of radio cocoons indicate that these objects are able to resists their very turbulent birth, it is reasonable to assume that they also survive the much weaker ambient turbulence. \\item A radio ghost can be revived to radio emission by the compression in a shock wave. This should lead to regions of diffuse synchrotron emission within clusters of galaxies, where accretion or merger shock waves are expected. Such objects are indeed observed, and are named cluster radio relics. Models for such processes can be found in En{\\ss}lin et al. \\cite*{1998AA...332..395E} and En{\\ss}lin \\& Krishna (2000)\\nocite{ensslin2000b}. \\item Another class of diffuse radio emitter in galaxy cluster might also be connected to ghosts: cluster radio halos. There is the speculative possibility that the short-lived radiating electrons are replenished by the decay of charged pions produced in hadronic collisions of long-lived cosmic ray protons, which were released from their confinement in radio ghosts \\cite{Ringberg99}.% \\item Although at low energies, the electron population within ghosts should remain mostly relativistic. Comptonization of the CMB photons by this electron population should produce CMB distortions, which have a spectral signature comparable to that of a grey absorber. This signature is weak with an optical depth of the order of $10^{-7}$ \\cite{ensslin2000a}, but principally observable. The energetically up-scattered CMB photons might also be detectable. \\end{itemize} For this work, to estimate the deflection of UHECR by ghosts, most of their above discussed properties are not of any importance. Only the magnetic fields are required, and their existence is nearly guaranteed by the existence of radio galaxies, and the extremely low magnetic diffusivity in extragalactic plasmas. Therefore, in the context of this work, the term `radio ghost' can be simply translated into `remnant magnetic fields of a former radio cocoon'. Even if turbulence is able to disrupt ghosts, the shear motions should increase the magnetic energy content of the fields, leading to an enhanced efficiency in UHECR deflection. We summaries, that the existence of regions in the IGM filled with remnant magnetic fields of former radio galaxies -- as it is assumed in this work -- is difficult to be doubted. ",
        "conclusions": "The, to some extent unexpected, isotropy in the UHECR flux observed particularly by Yakutsk and AGASA, is a very strong constraint to production and propagation models alike. Most of the scenarios proposed in the literature should produce a sizable anisotropy as either extragalactic luminous or dark matter is normally associated with the invoqued particle sources. If that is really the case, then it is our understanding of the topology and intensities of the intervening magnetic fields that is critically incomplete. The problem could reside inside the network of walls and filaments, in the interior of the large surrounding voids, or even in our {\\bf nearby} environment, namely the Galaxy halo.  In the present work we explore the possibility that radio ghosts, blobs of magnetized radio plasma remnant from past periods of activity in radio galaxies, being able to scatter UHECR in the intergalactic medium. Such a process could, in principle, degrade the direct incoming flux from the sources and build up a diffuse UHECR component large enough to be responsible for the observed degree of isotropy.  Our results show that, over the most conservative region of the radio ghost parameter space, such isotropization is not possible. This stands not from an inability of ghosts to scatter UHECR, but mainly from the fact that, under general conditions, ghosts should tend to cluster more strongly than the sources of the particles.  If, however, radio ghosts are able to buoy out into the surroundings of the dense large scale structures and into voids, while surviving the process, UHECR isotropy could be obtained in those cases in which the mean free path for interactions with ghosts is reduced below some few Mpc all over the propagation region ($\\sim b_{gh} < 0.5$).  Finally, one can question the basic assumption of this work, that radio ghosts are able to survive for cosmological times as intact objects. If radio ghosts mix with the ambient medium a significant magnetization of the IGM should result, which can be sufficiently strong in order to explain observed magnetic fields in clusters of galaxies. These magnetic fields would have short coherence length, and therefore do not violate Faraday rotation measurements. The ability of the fields to deflect UHECR particles would be increased compared to the case of intact, non-mixing ghosts for the following reasons: the average deflection of particles in a ghosts vicinity is proportional to $V_{\\gh}\\,B_{\\gh}$ (cross section $V_\\gh^{2/3}$ times scattering angle, the latter is proportional to $B_\\gh$ times $V_\\gh^{1/3}$). For a fixed magnetic energy budget $E_{B_\\gh} = V_{\\gh} \\,B_\\gh^2/(8\\pi)$ of the ghosts, this implies that the scattering efficiency scales with $V_{\\rm gh}^{1/2}$ and therefore increases with increasing size.  For fixed magnetic flux ($B_\\gh \\propto V^{-2/3}$) the scattering efficiency inreases with the size as $V_{\\rm gh}^{1/3}$.  This demonstrates that relaxing the hypothesis of compact ghosts makes this scenario even more powerful.\\\\[2em]   \\noindent {\\bf Acknowledgments}\\\\ \\noindent The authors wish to thank Dr. Simon White, Dr. Volker Springel, Dr. Avishai Dekel, Dr. Gerard Lemson and others for kindly making available the DM distribution from their large-scale-structure numerical simulation of the local Universe using a constrained realization of the initial perturbation field. It is a pleasure to acknowledge the constructive comments of the anonymous referee. GMT is partially supported by the Brazilian agencies FAPESP and CNPq. He thanks the MPA for the hospitality during his visit. TAE was temporary supported by the {National Science and Engineering Research Council of Canada} (NSERC)."
    },
    "0011/astro-ph0011348_arXiv.txt": {
        "abstract": "At low redshifts, a census of the baryons in all known reservoirs falls a factor of two to four below the total baryon density predicted from Big Bang nucleosynthesis arguments and observed light element ratios. Recent cosmological hydrodynamic simulations suggest that a significant fraction of these missing baryons could be in the form of warm/hot gas in the filaments and halos within which most field galaxies are embedded. With the release of source count results from Chandra and recent detections of this gas in \\Osix quasar absorption lines, it becomes interesting to examine the predictions and limits placed on this component of the X-ray background (XRB). We have used new hydrodynamical simulations to predict the total X-ray spectrum from the gas in the 100 eV to 10 keV range. We find that, when uncertainties in the normalization of the observed XRB and the value of $\\Omega_b$ are taken into account, our results are consistent with current observational limits placed on the contribution of emission from gas to the XRB. In the $0.5-2$ keV range, we expect the contribution from this component to be $0.63 \\times 10^{-12}$ \\cgs or between $6\\%$ and $18\\%$ of the extragalactic surface brightness. The peak fraction occurs in the $0.5-1$ keV range where the predicted line emission mirrors a spectral bump seen in the latest ASCA/ROSAT XRB data. ",
        "introduction": "Measurements of the total baryon density from the optical depth of the Lyman $\\alpha$ forest (Rauch et al. 1998) at high redshift ($z \\approx 2-3$) are in remarkable agreement with those derived from observed light-element ratios and nucleosynthesis arguments (BBN, Burles \\& Tytler), which yield a value of $\\Omega_{b,D/H} = 0.039 \\pm 0.002$, assuming $h \\equiv H_{0}/100 = 0.7.$ However, in the local universe, the combined observed contributions of stars, atomic and molecular hydrogen, and cluster gas fall a factor of two to four below this number (Fukugita, Hogan, \\& Peebles 1998). This is known as the missing baryon problem.   Cosmological hydrodynamical simulations suggest that, at low redshift, a significant fraction, up to $40\\%$, of all the baryons, could be in the warm/hot gas in filaments (e.g., Cen et al. 1995; Cen \\& Ostriker 1999; Dav\\'{e} et al. 2000; Croft et al. 2000). This diffuse gas, which is also known as the Warm/Hot Intergalactic Medium (WHIM), is located in filamentary structures between clusters of galaxies and in which groups of galaxies are embedded. It is shock-heated to temperatures $T = 10^{5}-10^{7}$K, between that of the hot cluster gas and the warm gas in voids. The comparatively low density and temperature of this gas make it a challenge to detect and as yet, attempts to observe it in emission have yielded marginal (Wang, Connolly \\& Brunner 1997; Boughn 1998; Kull \\& B\\\"ohringer 1999; Scharf et al. 2000) or negative results, with an upper bound of $1.1 \\times 10^{-12}$ \\cgs  in the 0.5 to 2 keV band for intercluster gas (Briel \\& Henry 1995).  However, warm gas has been detected as \\Osix absorption features in quasar spectra (Tripp \\& Savage 2000; Tripp, Savage, \\& Jenkins 2000, and references therein) which seem to coincide in redshift with galaxy overdensities. The WHIM may have been indirectly detected in the observed X-ray background (XRB) autocorrelation (Soltan \\& Freyberg 2000) and cross-correlation with the galaxy density distribution (Refregier, Helfand, \\& McMahon 1997; Soltan et al. 1997). In addition, rocket experiments show some evidence of \\Ose emission from this gas (D. McCammon et al., in preparation).  Most of the soft XRB has now been resolved into AGN and other sources, leading to an upper bound on the fraction of the extragalactic soft $1-2$ keV X-ray component due to the WHIM of $25\\%$ (Hasinger et al. 1998; Mushotzky et al. 2000; Giacconi et al. 2000). This is a small fraction of the XRB but could represent a large fraction of the baryons of moderate overdensity (Dav\\'{e} et al. 2000). Markevitch (1999), Phillips, Ostriker \\& Cen (2000) and Pierre, Bryan \\& Gastaud (2000) have already described possible methods to detect this gas using Chandra and XMM. In addition, the X-ray properties of this gas have been predicted from large scale simulations using both smoothed particle hydrodynamics (SPH) and Eulerian hydrodynamical methods (Croft et al. 2000; L. A. Phillips, in preparation).  We note that the hot gas component of the XRB is well established. Hot gas associated with rich clusters of galaxies is estimated to produce $\\sim 10\\%$ of the $1-2$ keV background (Hardy et al. 1998) and represents a fraction of approximately $8\\%$ of the computed total baryonic complement of the cosmic mass density (Fukugita et al 1998). These diffuse baryons are associated with the $\\sim 5\\%$ of all galaxies (Bahcall 2000) which are found in rich clusters. Most galaxies are in fact in lower density, lower velocity dispersion regions (such as poor clusters, groups, filaments and the field) and would be expected to be associated with lower density, lower temperature gas. It is the purpose of this paper to ascertain if the expected emission from this gas is consistent with our knowledge of the relatively soft XRB.  We present the total 100 eV to 10 keV X-ray spectrum from the WHIM calculated from new cosmological hydrodynamical simulations. The simulation is described in Section 2 and discussed in the context of the current limits placed on the different components of the XRB, including point sources, cluster gas and the more diffuse WHIM in Section 3. We summarize our results and briefly discuss the implied prospects of detecting the WHIM with ongoing and future X-ray missions in the Conclusions (Section 4). ",
        "conclusions": "Our prediction for the fractional X-ray contribution to the total XRB spectrum is consistent with the latest limits from AGN and other resolved sources in deep pointings from ROSAT (Hasinger et al. 1998) and Chandra (Mushotzky et al. 2000; Giacconi et al. 2000). The average integrated surface brightness from the gas of $0.63 \\times 10^{-12}$ \\cgs in the 0.5 - 2 keV band falls below the upper limits set by Briel \\& Henry (1995) and therefore, as of yet, the presence of gas between pairs of clusters cannot be ruled out. We are just reaching the flux limits necessary for the detection of this gas. The expected average flux levels lead us to predict that the WHIM could be observable with current missions such as Chandra and XMM. This possibility is explored in greater detail for XMM in Pierre et al. (2000) and for Chandra in Phillips et al. (2000). In the next decade, an effort should be made to quantify this WHIM component in clusters, groups, galaxies, filaments and/or sheets since a significant fraction of the baryons in the universe could be in this medium.  One of the major challenges in detecting the intermediate temperature gas in filaments will be to confirm the extragalactic origin of emission whose broad band spectral signature matches that predicted for the WHIM. The cross-correlation of the observed XRB and projected galaxy density distribution proves to be an ideal tool for this (L. A. Phillips, in preparation). The cross-correlation (CCF) flattens into a signature plateau at at scales of a few to ten arcmin (Refregier et al. 1997; Soltan et al. 1997). Soltan \\& Freyberg (2000) also see this feature in XRB autocorrelation calculations. The plateau in the CCF can be reproduced by the correlation of the X-ray emission from the WHIM in filaments with the simulated projected galaxy distribution (Phillips et al. 1999; Croft et al. 2000; L. A. Phillips, in preparation). If the gas is not detected at the predicted fluxes, the metallicity history of the WHIM in the simulations should be closely re-examined since a large fraction of the soft X-ray emission from this gas arises from metal line emission. Indeed, the peak fractional contribution from the WHIM occurs in the $0.5-1$ keV range where we predict a bump due to line emission coincident with a similar bump seen in the observed ASCA and ROSAT XRB spectrum.  The presence of the WHIM in groups of galaxies has already been established (Mulchaey et al. 1996). Groups that reside in galaxy filaments tend to have higher diffuse X-ray fluxes, indicating a higher density of gas (Mahdavi et al. 2000). Warm gas has also been detected in the halos of galaxies and in galaxy filaments (both in \\Osix absorption and in X-ray emission, see references above). Therefore the question is not whether there is a WHIM, but rather how much of it there is and where it is located. And once the answer to this question is known, we can move towards a better understanding of the physical state of most of the baryons in the universe."
    },
    "0011/astro-ph0011381_arXiv.txt": {
        "abstract": "We describe the implications of cosmic microwave background (CMB) observations and galaxy and cluster surveys of large scale structure (LSS) for theories of cosmic structure formation, especially emphasizing the recent Boomerang and Maxima CMB balloon experiments.  The inflation-based cosmic structure formation paradigm we have been operating with for two decades has never been in better shape. Here we primarily focus on a simplified inflation parameter set, $\\{\\omega_b,\\omega_{cdm},\\Omega_{tot}, \\Omega_\\Lambda,n_s,\\tau_C, \\sigma_8\\}$. Combining all of the current CMB+LSS data points to the remarkable conclusion that the local Hubble patch we can access has little mean curvature ($\\Omega_{tot} = 1.08\\pm 0.06$) and the initial fluctuations were nearly scale invariant ($n_s =1.03 \\pm 0.08$), both predictions of (non-baroque) inflation theory. The baryon density is found to be slightly larger than that preferred by independent Big Bang Nucleosynthesis estimates ($\\omega_b\\equiv \\Omega_b {\\rm h}^2 = 0.030\\pm 0.005$ cf. $0.019\\pm 0.002$). The CDM density is in the expected range ($\\omega_{cdm}=0.17 \\pm 0.02$). Even stranger is the CMB+LSS evidence that the density of the universe is dominated by unclustered energy akin to the cosmological constant ($\\Omega_\\Lambda = 0.66 \\pm 0.06$), at the same level as that inferred from high redshift supernova observations. We also sketch the CMB+LSS implications for massive neutrinos. \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011038_arXiv.txt": {
        "abstract": "{ We present the first high resolution photospheric X-ray spectrum of a Supersoft X-ray Source, the famous CAL~83 in the Large Magellanic Cloud. The spectrum was obtained with the Reflection Grating Spectrometer on {\\it XMM-Newton} during the Calibration/Performance Verification phase of the observatory. The spectrum covers the range 20--40 \\AA\\ at an approximately constant resolution of 0.05 \\AA, and shows very significant, intricate detail, that is very sensitive to the physical properties of the object. We present the results of an initial investigation of the spectrum, from which we draw the conclusion that the spectral structure is probably dominated by numerous absorption features due to transitions in the L-shells of the mid-$Z$ elements and the M-shell of Fe, in addition to a few strong K-shell features due to CNO. ",
        "introduction": "Supersoft X-ray Sources (or Supersoft Sources, SSS) were among the first new discoveries made with {\\it ROSAT}---in the first PSPC images of the LMC, the SSS stood out immediately (Tr\\\"umper et al. 1991). In retrospect, some of these objects had been seen before with {\\it Einstein}, but the {\\it ROSAT} detection in short order of several SSS focused attention on them as a class.  It is now generally accepted that the SSS are white dwarf stars undergoing steady nuclear burning in their envelopes. The inferred size of the radiating surface area, coupled with the characteristically very soft optically thick spectra, point to this idea as the natural interpretation for the observed emission. Some SSS are central stars of Planetary Nebulae, or related objects: shell-burning, (mostly) single stars on their way to the white dwarf graveyard. The classical SSS, however, is a high-luminosity ($L \\gax 0.1~L_{\\rm Edd}$), soft ($kT \\lax 50$ eV) low-mass binary with a white dwarf primary. In contrast to the case for neutron stars, for white dwarfs the efficiency of nuclear burning exceeds the efficiency of gravitational energy conversion by accretion, by a factor of order 100, so that SSS can sustain a high luminosity with a sub-Eddington accretion rate.  The basic model for the classical SSS has been worked out by van den Heuvel et al. (1992). A relatively high-mass white dwarf (0.7--1.2 $M_{\\sun}$) can sustain steady nuclear burning of accreted hydrogen near its surface if the accretion rate is in the range $1-4 \\times 10^{-7} M_{\\sun}$ yr$^{-1}$. Such an accretion rate can arise if the mass donor companion star is more massive than the white dwarf (1.5--2.0 $M_{\\sun}$). Mass transfer to the less massive star will cause the binary orbit to shrink, and the donor will transfer mass on a rapid thermal timescale. The predicted luminosities are in the range $6 \\times 10^{36} - 1 \\times 10^{38}$ erg s$^{-1}$, roughly in the observed range for SSS. The binary evolution aspects of this model, including an estimate for the birthrate and prevalence of SSS binaries, have been addressed by Rappaport et al. (1994).  An important part of the interest in the SSS stems from the fact that they may be a significant channel to SN Ia events. Mass is deposited onto the white dwarf at such a high rate that the accreted matter can burn non-eruptively, which makes the white dwarf mass grow (in contrast to low accretion-rate binaries, such as classical novae, which expel most of the previously accreted matter in explosive events), eventually reaching the Chandrasekhar mass. Rappaport et al. (1994) estimate that the rate of Galactic SN Ia supernovae associated with SSS binary evolution could be as high as $\\sim 0.006$ yr$^{-1}$.  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{XMM44_f1.ps}} \\caption{ Exposure corrected, background subtracted spectrum of CAL 83. We show the spectrum measured with RGS 1 (black histogram) overlaid on the spectrum measured with RGS 2 (red histogram), to emphasize the correspondence of coherent features in the spectrum in both spectrometer data sets. The wavelength bin equals 5 original RGS 81 $\\mu$m CCD spatial bins, which corresponds to approximately 0.06 \\AA, slightly larger than one resolution element. The range $20.1-23.9$ \\AA\\ is not present in the data from RGS2, because of malfunction of the drive electronics for one CCD chip in this spectrometer. Included for reference are wavelengths of the strongest resonance transitions in hydrogen- and helium-like C. } \\end{figure*}  A high resolution soft X-ray spectroscopic observation can in principle directly and definitively settle any remaining doubts about the nature of the emission and the compact object, by detecting characteristic photospheric structure in the spectrum. With a high resolution spectrum, we can attempt to obtain an accurate estimate for the fundamental stellar and binary parameters (luminosity, effective temperature, radius and mass of the white dwarf, mass transfer rate), which has a direct bearing on the evolutionary status and  ultimate fate of the binary. In this {\\it Letter}, we present preliminary results of a high resolution spectroscopic observation of the important SSS CAL 83 (Cowley et al. 1984, Pakull et al. 1985, Crampton et al. 1987, Smale et al. 1988, Parmar et al. 1998) which directly addresses the above issues. ",
        "conclusions": "We have presented the first high resolution X-ray spectrum of CAL 83. With {\\it XMM-Newton} RGS, we finally conclusively demonstrate the photospheric origin of the soft X-ray emission. Moreover, our spectrum qualitatively resembles the spectrum expected for the emission from a very hot white dwarf star, and therefore confirms the basic model for binary Supersoft X-ray Sources.  We detect radiation from the photosphere between 20 and $\\sim 40$ \\AA, at an approximately constant wavelength resolution of 0.05 \\AA. The spectrum shows very significant absorption structure. From comparison to model atmospheres calculations at $kT_{\\rm eff} \\sim 45$ eV, $\\log g \\sim 8.5$, for varying abundances, we conclude that most of this structure is due to the superposition of numerous, unresolved narrow absorption features due to transitions in the L shells of the mid-$Z$ elements, the M shell of Fe, as well as in the K shell of C. A quantitatively reliable estimate of the stellar parameters has to await a detailed analysis in terms of a series of dedicated model atmospheres calculations."
    },
    "0011/astro-ph0011512_arXiv.txt": {
        "abstract": "We present observations of \\c2h5oh toward molecular clouds in Sgr A, Sgr B2 and associated with thermal and non-thermal features in the Galactic center. \\c2h5oh emission in Sgr A and  Sgr B2 is widespread, but not uniform. \\c2h5oh emission is much weaker or it is not detected in some molecular clouds in both complexes, in particular those with radial velocities between 70 and 120 \\kms. While most of the clouds associated with the thermal features do not show \\c2h5oh emission, that associated with the Non-Thermal Radio Arc shows emission. The fractional abundance  of \\c2h5oh in most of the clouds with radial velocities between 0 and 70 \\kms in Sgr A and Sgr B2 is relatively high, of few \\E{-8}. The \\c2h5oh abundance decreases by more than one order of magnitude (\\lsim \\E{-9}) in the clouds associated with the thermal features. The large abundance of \\c2h5oh in the gas-phase indicates that \\c2h5oh has formed in grains and released to gas-phase by shocks in the last \\apro \\E{5} years. The implications of this finding in the origin of the shocks in the GC is briefly discussed. ",
        "introduction": "Ethanol (\\c2h5oh) was first detected by \\citet{zuckerman75} towards the massive star forming region in the Galactic center (GC) Sgr B2. Subsequent observations of this molecule in the interstellar medium have shown that \\c2h5oh  is only present in the dense (\\gsim\\E{6}\\cmm3) and hot (\\gsim 100K) cores associated with newly formed massive stars \\citep{irvine87,millar88,ohishi95,millar95,nummelin98}. In hot cores, the abundance of \\c2h5oh in gas-phase ranges from \\E{-8} to \\E{-9}, and gas-phase chemistry cannot account even for the lowest fractional abundance of this molecule. Therefore, \\c2h5oh emission is considered one of the best tracers of dust chemistry \\citep{millar88,charnley92,charnley95}. Furthermore, the transient nature of the alcohol chemistry in molecular clouds, makes the abundance of \\c2h5oh an excellent clock to estimate when this molecule was injected to the gas-phase \\citep{charnley95}.  The GC molecular clouds show high kinetic temperatures (\\gsim 80 K) \\citep{hutte93, rodriguez01} and  widespread large abundance of refractory molecules like SiO \\citep{martin97}. Shock waves have been invoked to explain the heating, the morphology and the relatively large abundance of SiO in the GC \\citep{martin97}. Shocks are also expected to sputter molecules from the icy grain mantles increasing the gas-phase abundance of the molecules formed in grains. Evidences for the release of molecules from grain mantles in selected molecular clouds close to Sgr A comes from the detection of \\c2h5oh and HOCO$^+$\\ in this region \\citep{minh92, charnley00}. In the scenario that shocks drive the chemistry in the GC clouds, it is expected that widespread \\c2h5oh emission closely follows that of SiO.  In this letter we present observations of \\c2h5oh toward molecular clouds in Sgr A, Sgr B2 complexes and associated with the thermal features (Arched Filaments --hereafter TAF-- and the Sickle) and the Non Thermal Radio Arc (NTRA) in the GC \\citep[for the nomenclature]{lang99}. We find that the \\c2h5oh emission in the GC is widespread, but large variations in the abundance of this molecule are found between the different GC molecular clouds. The large \\c2h5oh abundance and its variations are consistent with a scenario in which recent (\\lsim \\E{5} years) shocks dominate the chemistry in the GC by grain processing.   \\section {Observations and results}  The observations of the \\ct, \\sc\\ and \\no\\ lines of \\c2h5oh towards the GC were carried out simultaneously with the IRAM 30-m telescope. To estimate the \\hdos\\ column density we also observed simultaneously the \\uc\\ line of \\tco\\ and \\cdo, the \\du\\ line of \\cdo\\, and the \\td\\ line of CS. The half power beam width of the telescope was 24\\as, 17\\as\\ and 12\\as\\ for the 3, 2 and 1.3 mm bands. The receivers, equipped with SIS mixers, were tuned to single side band with image rejections of \\gsim 10 dB. The typical system temperatures were 300, 500 and 900 K for the 3, 2 and 1.3 mm lines respectively. We used two filter banks of 256$\\times$1MHz and one of 512$\\times$1MHz as spectrometers. The velocity resolution provided by the filter banks were 3, 2 and 1.3 \\kms for the 3, 2 and 1.3 mm bands respectively. The calibration was achieved by observing hot and cold loads. The line intensities are given in units of \\ta*.  The observed molecular clouds were selected from the SiO maps of \\cite{martin97} and their locations are shown in the upper panel of Fig. 1. The typical line profiles of \\c2h5oh, \\tco\\ and CS towards selected clouds in SgrA (M-0.11-0.08), in Sgr B2 (M+0.76-0.05), and associated with the Sickle (M+0.20-0.03) and the NTRA (M+0.17+0.01) are shown in the lower panels of Fig. 1. Table 1 summarizes the results obtained for all molecular clouds. The emission in the \\ct\\ line of \\c2h5oh is widespread in Sgr A and Sgr B2. However, the line profiles of \\c2h5oh are different from those of CS, \\tco, and \\cdo. This is illustrated in Fig. 1 for M+0.76-0.06. The CS and the \\tco\\ emission appears from -20 \\kms to 120 \\kms. However, like the SiO emission in Sgr A and Sgr B2 \\citep{martin00}, the \\c2h5oh emission mainly appears for radial velocities between 10 and 70 \\kms. This indicates variations in the abundance of \\c2h5oh in the molecular clouds along the line of sight  with different radial velocities.  Column 4 in Table 1 gives the radial velocity of the peak intensity of the \\ct\\ \\c2h5oh line and the velocity range where the \\c2h5oh emission is  detected.  The morphology of the \\c2h5oh emission in the GC is also different from those of CS, \\tco, and \\cdo. \\c2h5oh emission is not detected towards M-0.08-0.06 in SgrA. Furthermore, the \\c2h5oh emission shows different behavior towards the different kind of filaments in the GC. In general, the molecular clouds associated with the thermal features like the Sickle  \\citep[M+0.18-0.04 and M+0.20-0.03]{serabyn91} and the TAFs E1 and E2 \\citep[M+0.13+0.02 and M+0.17+0.01]{serabyn87} do not show \\c2h5oh emission. An interesting exception is the detection of \\c2h5oh emission towards M+0.04+0.03, a molecular cloud located south of the TAF W1. However, the rather narrow radial velocity range of the \\c2h5oh emission as compared with those of the CS and \\tco \\ emission \\citep{serabyn87} suggests that the bulk of the molecular gas associated with thermal features does not show \\c2h5oh emission. Mapping of the \\c2h5oh emission in this region will tell the possible association of this velocity component with the thermal features. In contrast to the lack of \\c2h5oh emission from the molecular clouds associated with the thermal features, weak \\c2h5oh emission is detected towards M+0.17+0.01 in the NTRA.  In summary, \\c2h5oh emission in the GC molecular clouds is widespread, but it shows substantial differences between the different features in the GC. Furthermore, the \\c2h5oh emission is mainly restricted to molecular clouds with radial velocities between 10 and 70 \\kms. ",
        "conclusions": ""
    },
    "0011/astro-ph0011128_arXiv.txt": {
        "abstract": "Specially-designed microlensing searches, some of which have been underway for several years, are sensitive to extrasolar planets orbiting the most common stars in our Galaxy. Microlensing is particularly well-suited to the detection of Jupiter-mass planets orbiting their parent stars at several AU. Since Jovian analogs are thought to influence the subsequent evolution of most planetary systems, they are particularly important to study. The orbital radii and distances to the planetary systems probed by microlensing are larger than those currently studied by radial velocity techniques; the two methods are thus complementary. Recent results from microlensing searches are discussed, including constraints on Jovian analogs orbiting typical Galactic stars.  Benefits and drawbacks of the technique for the characterization of planetary systems, and future prospects are briefly reviewed. ",
        "introduction": "The first detection of an extrasolar planet around a normal star (Mayor \\& Queloz 1995) and subsequent deluge of similar discoveries by the radial velocity technique (Marcy, Cochran \\& Mayor 2000), have taught us that $\\sim$5\\% of solar type stars harbor planetary systems very {\\it unlike\\/} our own. What remains to be determined is the abundance of planetary systems {\\it similar\\/} to that of the Sun's (eg., terrestrial planets at $\\sim$0.5-2~AU or Jovian analogs at $\\sim$3-12~AU) and the frequency of planets around the most common stars in our Galaxy, M dwarfs. Microlensing is providing the first partial answers to these questions. ",
        "conclusions": ""
    },
    "0011/astro-ph0011402_arXiv.txt": {
        "abstract": "We present synthetic Doppler maps of gaseous flows in semidetached binaries based on the results of 3D gas dynamical simulations. Using of gas dynamical calculations alongside with Doppler tomography technique permits to identify main features of the flow on the Doppler maps without solution of ill-posed inverse problem. Comparison of synthetic tomograms with observations makes possible both to refine the gas dynamical model and to interpret the observational data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011544_arXiv.txt": {
        "abstract": "The recent Boomerang and MAXIMA data on the cosmic microwave background suggest a large value for the baryonic matter density of the universe, $\\omega_b \\sim 0.03$.  This density is larger than allowed by standard big bang nucleosynthesis theory and observations on the abundances of the light elements. We explore here the possibility of accommodating this high density in inhomogeneous big bang nucleosynthesis (IBBN).  It turns out that in IBBN the observed $\\D$ and $Y_p$ values are quite consistent with this high density.  However, IBBN is not able to reduce the $\\ZLi$ yield by more than about a factor of two. For IBBN to be the solution, one has to accept that the $\\ZLi$ plateau in population II halo stars is depleted from the primordial abundance by at least a factor of two. ",
        "introduction": "The most accurate way to estimate the average density of baryonic matter in the universe, $\\rho_b$, has for a long time been big bang nucleosynthesis. In standard big bang nucleosynthesis (SBBN)\\cite{SBBN} the calculated primordial yields of the light elements depend only on the baryon-to-photon ratio $\\eta \\equiv n_b/n_\\gamma$. Future precision measurements of the fluctuations in the cosmic microwave background (CMB) will provide another way of measuring the baryon density.  The shape of the angular power spectrum of these fluctuations will depend on a number of cosmological parameters, among which is the baryon-to-photon ratio.  In this context it is usually given as the baryonic contribution to the critical density times the Hubble constant squared, $\\omega_b \\equiv \\Omega_bh^2$.  These two measurements rely on completely different physics, and if they agree it is an important confirmation that we have the right understanding of the early universe.  Assuming that there was no significant entropy production\\cite{KT00} in the universe between BBN and recombination, the two parameters $\\omega_b$ and $\\eta$ are related via the present temperature of the CMB, $T = 2.725$ K, by \\beq \\et = 274 \\omega_b. \\eeq  Two recent balloon-borne experiments, Boomerang\\cite{boom1} and Maxima-1\\cite{max1}, have now provided us with the first measurements of the CMB angular power spectrum which are of sufficient quality that an estimate of $\\omega_b$ can be made from them. The result\\cite{boommax}, $\\omega_b \\sim 0.030$, is significantly higher than the SBBN result.  Indeed, SBBN clearly cannot accommodate as high a baryon density as $\\omega_b = 0.030$\\cite{BoomSBBN}.  The SBBN yields for $\\omega_b = 0.030$ are $Y_p = 0.251$, $\\DH = 1.7\\times10^{-5}$, $\\EHeH = 8.4\\times10^{-6}$, and $\\ZLiH = 7.6\\times10^{-10}$.  With the exception of $\\ZLi$, the uncertainty in these SBBN yields is much less than the uncertainty in the primordial abundances derived from observations.  There is much debate about chemical evolution and systematic effects in observations. For the primordial helium abundance there are two competing estimates, the ``low $\\UHe$''\\cite{loHe}, $Y_p = 0.234\\pm0.003$ and the ``high $\\UHe$''\\cite{hiHe}, $Y_p = 0.244\\pm0.002$.  The difference is largely due to the method of estimating the present abundances from the observed line intensities, suggesting that the systematic uncertainty may be larger than the $\\pm0.005$ usually assumed\\cite{Olive00}.  Burles and Tytler \\cite{Tytler} claim to have established the primordial deuterium abundance  as $\\DH = (3.3\\pm0.25)\\times10^{-5}$ (``low D\") based on Lyman-series absorption by high-redshift clouds.  This is based on a detected low deuterium abundance in three such systems and upper limits from others.  There remains one such system, where a high deuterium abundance, $\\DH \\sim 2\\times10^{-4}$ (``high D\"), has been observed\\cite{hiD}.  It may be that the accuracy of such observations has been overestimated\\cite{mesoturb}. We refer the reader to recent reviews\\cite{BBNreviews} for further discussion and adopt the observational constraints $Y_p = 0.228$--$0.248$, $\\DH = 2.9$--$4.0\\times10^{-5}$\\cite{Steigman}.  These lead to the SBBN range $\\omega_b = 0.004$--$0.021$ from $Y_p$ and $\\omega_b = 0.018$--$0.022$ from $\\DH$. The ``low D'' of Burles and Tytler gives $\\omega_b = 0.020\\pm0.01$.  The ``Spite plateau''\\cite{Spite} of $\\ZLi$ abundance in population II halo stars provides us with the best estimate of the primordial lithium abundance.  Bonifacio and Molaro\\cite{BoMo97} obtain a present lithium abundance $\\lgZLiH = -9.80\\pm0.012\\pm0.05$ for these stars, and argue against any significant depletion from the primordial abundance, based on the lack of dispersion in the data. Pinsonneault \\etal\\cite{Pins} estimate a depletion factor of 0.2--0.4 dex from rotationally induced mixing. In a more recent study Ryan \\etal\\cite{Ryan99} obtain a mean abundance $\\lgZLiH = -9.88$ for the Spite plateau, and argue that the narrow spread, less than 0.02 dex, limits depletion by rotationally induced mixing to less than 0.1 dex. Moreover, they observe a slight trend with metallicity, suggesting a galactic contribution, so that the primordial abundance could be lower than the observed abundance. Based on this, Ryan \\etal\\cite{Ryan00} derive the range $\\lgZLiH = -10.04$\\ldots$-9.72$ for the primordial abundance. This very tight upper limit to primordial lithium is in conflict with the ``low D'' estimate\\cite{Tytler}, $\\omega_b = 0.020$, which corresponds to $\\lgZLiH = -9.5$ in SBBN.  At present there seems to be a consensus\\cite{Pins,VC,Deli}, that the primordial abundance cannot be larger than the observed value by more than about 0.4 dex, so that models producing a primordial lithium abundance above $\\lgZLiH = -9.4$, or $\\ZLiH = 4\\times10^{-10}$, would be ruled out.  This gives an upper limit $\\omega_b \\leq 0.022$ in SBBN.  While the upper limit to $\\omega_b$ from $Y_p$ is rather soft, the deuterium and lithium abundances are in clear conflict with $\\omega_b = 0.03$.  This CMB measurement of $\\omega_b$ is rather preliminary and not very accurate, and thus it would be premature to discard the SBBN estimate.  Since the CMB measurement is a simultaneous fit of many cosmological parameters, the lack of prominence of the second peak in the power spectrum could have another explanation.  Suggestions include a significant tilt\\cite{TZboomWSP00}, an isocurvature component\\cite{EKV00}, or a feature\\cite{Barriga} in the primordial perturbation spectrum, and topological defects\\cite{topo}.  Future satellite experiments, MAP\\cite{MAP} and Planck\\cite{Planck}, will give us a better CMB determination of $\\omega_b$.  If the situation persists and the CMB value of $\\omega_b$ turns out to be $\\omega_b \\sim 0.03$ or higher, one has to consider abandoning SBBN in favor of non-standard BBN\\cite{NSBBN}. There are many modifications suggested to standard BBN, some of which can accommodate a higher baryon density\\cite{NSBBNhighden}.  We consider here one such possible modification, inhomogeneous big bang nucleosynthesis (IBBN).  In IBBN one assumes that the baryon density was inhomogeneous during BBN.  This inhomogeneity could be an initial condition resulting from unknown physics in the early universe, or it could be the result of some physical event, like a phase transition, occurring before BBN.  The nature of IBBN depends on the distance scale of the inhomogeneity.  Especially interesting is the case, where this distance scale is of the order of the neutron diffusion length during BBN\\cite{earlyIBBN}.  Then neutron diffusion will lead to an inhomogeneity in the neutron-to-proton ratio.  If the density contrast between high- and low-density regions is large enough, this ratio can exceed unity in the low-density region.  In this case IBBN leads to a decrease in the $\\UHe$ yield and an increase in the $\\D$ yield, favoring larger $\\omega_b$.  However, the $\\ZLi$ yield is more likely to increase than decrease from the SBBN result for the same average $\\omega_b$.  In this paper we focus on the question, whether the high value $\\omega_b = 0.030$, corresponding to $\\et = 8.22$, would be acceptable in IBBN.  There have been many papers on IBBN\\cite{earlyIBBN}. However, because of the large number of parameters in the IBBN model, the parameter space has not been covered thoroughly in them.  Focusing on one particular value of $\\et$ we reduce the dimension of the parameter space, allowing us to search the space of the remaining parameters more completely.  We can thus give a more specific answer to the above question. ",
        "conclusions": "It is intriguing, that the ``high'' baryon density $\\omega_b \\sim 0.03$, suggested by recent CMB results fits so well with the ``low'' $\\DH$ and ``high'' $Y_p$ observations in the IBBN scenario. However, the $\\ZLi$ yield is significantly higher than the standard estimates for its primordial value.  For the SS geometry, IBBN can reduce the lithium yield by about a factor of two in the best case, but for $\\omega_b = 0.03$, we still get $\\ZLiH = 4\\times10^{-10}$, requiring depletion by at least a factor of two in the Spite plateau. This may be marginally acceptable, especially if one allows for the uncertainty in the $\\ZLi$ yield due to the reaction rates.  We note for comparison, that the widely accepted ``low-D\" determination $\\omega_b = 0.020$ in SBBN yields $\\ZLiH = 3.4\\times10^{-10}$ and thus also requires significant lithium depletion\\cite{BNT00}.  The better developed scenarios for producing the kind of inhomogeneity needed for IBBN tend to lead to the CC type of geometry.  For this geometry, IBBN is not able to reduce the lithium abundance significantly below the SBBN result, so that depletion by a factor of four would be needed. This seems to be considered unacceptable at present."
    },
    "0011/astro-ph0011258_arXiv.txt": {
        "abstract": "We examine the effects of galaxy outflows on the formation of dwarf galaxies in numerical simulations of the high-redshift Universe. Using a Smoothed Particle Hydrodynamic code, we conduct two detailed simulations of a (5.2 Mpc/$h$)$^3$ comoving volume of the Universe. In both simulations we implement simple, well-motivated models of galaxy identification and star formation, while our second simulation also includes a simple ``blow-out'' model of galaxy outflows in which supernova driven winds from newly formed disk galaxies punch-out and shock the intergalactic medium while leaving the host galaxies intact. A direct comparison between these simulations suggests that there are two major mechanisms by which outflows affect dwarf formation. Firstly, the formation of an outflow slows down the further accretion of gas onto a galaxy, causing an overall decrease of approximately 50\\% in the total gas mass accreted by the objects in our simulations.  Additionally, our simulations uncover a significant population of $\\sim 10^9 M_\\odot$ objects whose formation is suppressed by the mechanism of ``baryonic stripping,'' in which outflows from early galaxies strip the gas out of nearby overdense regions that would have otherwise later formed into dwarf galaxies. This mechanism may be important in explaining the observed discrepancy between the number of dwarf galaxies predicted and observed in the local group and provide a natural explanation for the formation of empty halos which may be required by the existence of the extremely gas-poor extra-galactic High-Velocity Clouds. ",
        "introduction": "It has long been recognized that the properties of the diffuse gas within clusters of galaxies are indicative of a violent past. The slope of the X-ray luminosity-temperature relationship of the hot intracluster medium  in galaxy clusters is far too steep to be  due to heating only by viralising shocks, requiring that large quantities of heated material be injected into this gas at early times. (see eg.\\ Kaiser 1991; Mushotzky \\& Scharf 1997; Eke, Navarro \\& Frenk 1998; Cavaliere, Menci, \\& Tozzi 1999). Similarly, the metallicity of the intracluster medium is observed to be quite high, $\\sim .3 Z_\\odot$, and constant over a large range of cluster masses (Renzini 1997), indicating widespread enrichment by material ejected by supernovae.  (For an alternative viewpoint, see however Bryan 2000 and Pearce et al.\\ 2000).   While the first observational evidence for preheating and enrichment was met with theoretical resistance, hindsight leads us to believe that we might have expected this all along.  As structure formation is thought to occur hierarchically, and this growth is biased to the same over-dense regions at all times, it is natural to expect that a large population of dwarf galaxies would have once been found near the locations of present-day clusters. It has also been widely observed that dwarf galaxies undergo outflows both at low and high redshift.  Several studies of expanding HI gas in nearby dwarf galaxies show clear evidence of dense, expanding shells with velocities above 15 km/s (Marlowe et al.\\ 1995; Heckman 1997; Hunter et al.\\ 1998; Martin 1998) as well as halos of hot ($\\sim 5 \\times 10^6$ K) gas surrounding these objects (della Ceca et al.\\ 1996; Bomans, Chu, \\& Hopp 1997). Similarly, spectroscopic studies of galaxies have confirmed that high-velocity winds are present around dwarf galaxies at redshifts $\\sim$ 3 (Pettini et al.\\ 1998; Pettini et al.\\ 2001) and even higher (Frye \\& Broadhurst 1998; Warren et al.\\ 1998). Taken together these observational facts suggest a picture in which outflows powered by starbursts in the earliest dwarf galaxies had a huge impact on the properties of the gas that later condensed within clusters of galaxies.  Thus a vague theoretical scheme exists for understanding the preheating of clusters; yet surprisingly, the impact of this mechanism on the intergalactic medium (IGM) has been largely unexplored. Although clusters and galaxies both condense out of the same material, preheating is only considered in cluster simulations (see eg.\\ Metzler \\& Evrard 1994; Yepes at al.\\ 1997) and simulations of the properties of the IGM (Nath \\& Trentham 1997; Madau, Ferrara, \\& Rees 2000), while all numerical galaxy formation simulations carried out to date assume primordial conditions. Feedback in these simulations is modeled solely as impacting the interstellar medium  within the galaxy, without affecting the IGM and neighboring galaxy formation.  We have recently conducted two exploratory semi-analytical studies on the impact of galaxy outflows on the IGM and the formation of galaxies within it (Scannapieco, Ferrara, \\& Broadhurst 2000; Scannapieco \\& Broadhurst 2001, hereafter SB).  These studies have shown that outflows may be crucial in resolving such long-standing astronomical mysteries as the factor of four discrepancy between the number of observed and predicted Milky-Way Satellite galaxies (see eg.\\ Klypin et al.\\ 1999; Moore et al.\\ 1999) and the physical processes that cause elliptical galaxies to be typically much larger than disks (see eg.\\ Bromley et al.\\ 1998). The results of our semi-analytic simulations hinged on the importance of a mechanism of ``baryonic stripping'' in which outflows from early galaxies strip the gas out of nearby overdense regions of the IGM that would have otherwise later formed into dwarf galaxies.  In SB we were able to illustrate the importance of this mechanism in an idealized context, however the nonlinear clustering of galaxies, complicated evolution of outflows, and details of shock motion through collapsing regions can be captured only through numerical simulation.  In this article we address this question in the more detailed setting of Smoothed Particle Hydrodynamic (SPH) simulations, studying the impact of a simple model of galaxy outflows on the formation of structure in a typical volume of the Universe.   Feedback has been implemented by other authors using a number of different prescriptions. The lack of a fundamental theory of star formation and the interstellar medium render it extremely difficult to decide how to model star formation and feedback `correctly'. For example, the dispersal of energy from supernovae can be achieved by thermal heating, by passing kinetic energy to the gas, or by some arbitrary combination of the two. Katz (1992) was the first to consider the effect of thermal heating, and this seminal work highlighted a fundamental problem, namely that the cooling time for hot dense gas is so short (less than Myr) that simply returning feedback energy as thermal heating has little effect. Navarro \\& White (1993) considered a feedback model with both kinetic boosts and thermal heating. Mihos and Hernquist (1994) and Gerritsen (1997) have studied the effect of feedback modeled as kinetic boosts in isolated disk galaxies.  Gerritsen has argued that the boost model is too efficient when energy transfer parameters are chosen to match theoretical predictions such as those of Thorton et al.\\ (1998). This is probably related to the fact that the kinetic energy boosts are calculated using a fixed energy budget that does not account for the force softening that occurs at the short scales on which feedback is modeled.  More recently, Yepes et al.\\ (1997) have constructed a multi-phase model of star formation and feedback, which has been implemented in SPH simulations by Hultman \\& Pharasyn (1999). Given the large number of uncertainties in modeling the evolution of the interstellar medium, it remains unclear if such sophisticated approaches are well-motivated.  Thacker \\& Couchman (2000) present a number of different approaches to modeling feedback, including a new approach where drastic radiative losses are prevented by using a modified cooling formalism.  Springel (2000) has developed a model where the local evolution of gas can become dominated by turbulent pressure-support from exploding supernovae. A survey of the current status of feedback in simulations of galaxy formation can be found in Thacker \\& Davis (1999).  Given the inherent difficulty in modeling feedback, in this work we adopt a simple kinetic model that is inspired by the numerical simulations conducted by Mac Low \\& Ferrara (1999) and Ferrara \\& Tolstoy (1999).  Here the authors demonstrate that in disk galaxies within halos of $10^7 M_\\odot \\lesssim M \\lesssim 10^9 M_\\odot$, a ``blow-out'' occurs in which super-bubbles formed by groups of Type II supernovae punch out of the disk, shocking the IGM while leaving the interstellar medium of the galaxy intact.  This is due to the fact that while supernovae are inefficient at ejecting material within the disk, they expand almost freely in the perpendicular direction, resulting in a large organized wind made up of mostly the gas surrounding the galaxy.  This scenario has also been supported by the numerical simulations of Martel \\& Shapiro (2000).  The structure of this work is as follows.  In \\S2 we outline the numerical method adopted and the simplifying assumptions used in constructing our model for galaxy outflows.  In \\S3 we compare the results from hydrodynamic simulations conducted both with and without outflows to examine the impact of outflows on the number density and distribution of galaxies of varying mass scales.  Conclusions are given in \\S4. ",
        "conclusions": "Despite the overwhelming observational evidence for galaxy outflows at high and low redshift, the impact of this process on galaxy formation has been little explored.  While all numerical galaxy formation simulations carried out to date assume formation in primordial conditions, exploratory semi-analytical studies have suggested that many galaxies form under much different circumstances.  In this work, we have implemented a simple ``blow-out'' model for feedback from supernova generated winds, to study the impact of outflows on the formation of high-redshift dwarf galaxies in a characteristic volume of the Universe.  While these simulations avoid galaxy self-annihilation or ``blow-away'' by construction, they have uncovered two major mechanism by which outflows impact the formation of dwarf galaxies.  Firstly, outflowing dwarf galaxies impact their own formation by slowing down the further accretion of gas. This causes an overall decrease of about 50\\% in the total gas mass accreted by these objects which is roughly constant over the mass range studied in our simulations.  In addition to this mechanism, and perhaps more interestingly, our SPH simulations have also uncovered a significant number of would-be dwarves whose formation has been suppressed by baryonic stripping by neighboring objects.  In these cases, the winds from nearby objects reach the overdense region sufficiently early to eject the gas from the collapsing dark-matter perturbation.   This results in a large fraction of $\\sim 10^9 M_\\odot$ empty halos in which the dark matter has virialized but no baryonic gas has collapsed.  While our simulations are only able to identify this suppression at high redshift, this mechanism is suggestive of the discrepancy between the number of observed Milky-Way satellites and predictions from CDM models that do not include outflows (Klypin et al.\\ 1999; Moore et al.\\ 1999).  The dark-halos generated in this scheme also provide a natural explanation for the existence of massive yet extremely gas-poor extra-galactic High-Velocity Clouds (see eg.\\ Blitz et al.\\ 1999), as well as the observed lack of low-mass galaxies in the field relative to dense clusters (eg.\\ Phillips et al.\\ 1998). Note that as pointed out by Trentham, M\\\"oller, and Ramirez-Ruiz (2000) the existence of a large number of empty halos does not depend on the detailed predictions of the CDM model but follows simply from the assumption that the mass function of dark matter halos does not vary with environment, thus suggesting that the observed enviromentally dependent suppression of dwarf formation is due to astrophysical considerations rather than some unknown properties of the dark matter or the primordial power spectrum (eg.\\ Kamionkowski \\& Liddle 2000; Spergel \\& Steinhardt 2001).  Our results rely on a simple model of outflow formation in a limited volume of the Universe at redshifts above 6, and thus can not be directly compared with observations. Nevertheless the lesson to be drawn from these simulations is clear. Just as in the case of galaxy clusters, outflows  profoundly affect the formation and distribution of dwarf galaxies. These effects are not confined simply to the properties of the outflowing galaxies themselves but the formation of objects nearby, and this interaction between neighbors plays an important role in  the history of the formation of galaxies. While the assumption of primordial conditions is a good first step in the study of this problem then, we will never fully understand galaxy  formation until we have first understood mechanical feedback from galaxy outflows."
    },
    "0011/astro-ph0011291_arXiv.txt": {
        "abstract": "We present the results of HI mapping with the NRAO\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under contract with the National Science Foundation.} VLA of one of the most metal--deficient blue compact dwarf (BCD) galaxies known, SBS~0335--052, with an oxygen abundance of only 1/40 that of the Sun.  We study the structure and dynamics of the neutral gas in this chemically young object with a spatial resolution of 20\\farcs5 $\\times$ 15\\arcsec\\ ($\\sim$ 5.4 $\\times$ 3.9 kpc at an assumed distance of 54.3\\,Mpc), a sensitivity at the 2$\\sigma$ detection level of $\\sim$2.0 K or $ 7.5 \\times 10^{19}$ cm$^{-2}$ and a velocity resolution of 21.2 km\\,s$^{-1}$. We detected a large HI complex associated with this object with an overall size of about 66 by 22 kpc and elongated in the East--West direction.  There are two prominent, slightly resolved peaks visible in the integrated HI map, separated in the East--West direction by 22 kpc (84\\arcsec).  The eastern peak is nearly coincident with the position of the optical galaxy SBS~0335--052.  The western peak is about a factor of 1.3 brighter in the HI line and is identified with a faint blue compact dwarf galaxy, SBS 0335--052W, with $m_B =$ 19.4, and a metallicity close to the lowest values known for BCDs, about 1/50 that of the Sun. The radial velocities of both systems are similar, suggesting that the two BCDs SBS~0335--052 and SBS~0335--052W constitute a pair of dwarf galaxies embedded in a common HI envelope. Alternatively, the BCDs can be the nuclei of two distinct interacting primordial HI clouds.  The estimated total dynamical mass, assuming the BCDs form a bound system, is larger than $\\sim 6 \\times 10^{9}$\\,$M_{\\odot}$. This is to be compared to a total gaseous mass $M_{gas} = 2.1 \\times 10^{9}$ $M_{\\odot}$, and a total stellar mass $M_{star}$   $\\leq 10^{8}$ $M_{\\odot}$. Hence, the mass of the SBS~0335--052 system is dominated by dark matter. Because of the disturbed HI velocity field and the presence of what might be tidal tails at either end of the system, we favor the hypothesis of tidal triggering of the star formation in this system. It can be due to either the nearby giant galaxy NGC~1376 or the mutual gravitational interaction of the two HI clouds. ",
        "introduction": "Galaxy formation remains a key issue in observational cosmology. While much progress has been made in finding large populations of galaxies at high ($z > 3$) redshifts (e.g. Steidel et al. 1996; Dey et al. 1998), truly young galaxies in the process of forming, defined as objects which are experiencing their first epoch of star formation, remain elusive. The spectra of those distant galaxies generally indicate the presence of a substantial amount of heavy elements, implying previous star formation and metal enrichment.  We adopt here a different approach for studying truly young galaxies.  Instead of searching for high redshift objects, we look for nearby young dwarf galaxies which are possibly undergoing their first burst of star formation.  As pointed out more than a quarter of a century ago by Sargent \\& Searle (1970) and Searle \\& Sargent (1972), some low--mass galaxies in the local Universe may approximate galaxies in an early stage of their formation. Later, Thuan \\& Martin (1981) identified blue compact dwarf (BCD) galaxies as a class of low-luminosity ($M_B$ $\\ga$ --18) extragalactic systems undergoing a strong burst of star formation (SF) while exhibiting extremely low heavy element abundances.  For many years the best candidate for a truly young galaxy has been I~Zw~18 $\\equiv$ Mkn~116, with a metallicity $Z$ of about 1/50 $Z_{\\odot}$ (e.g. Lequeux et al. 1979). Studies of this galaxy in the last decade suggest indeed that it is chemically unevolved. Its extremely low metallicity derived from the HII region-like spectrum has been confirmed by many subsequent studies (see Izotov \\& Thuan (1998) and references therein).  Observations in neutral hydrogen (HI) show that it is a very gas--rich system (Lequeux \\& Viallefond 1980; Viallefond et al.\\ 1987; van Zee et al.\\ 1998a). Also, its HI gas appears to have an extremely low metallicity (Kunth et al.\\ 1994), although this last point is considered controversial (Pettini \\& Lipman 1995; van Zee et al. 1998a).  Until the end of the eighties, I~Zw~18 remained the only BCD with a heavy element abundance as low as $Z_\\odot$/50. Despite efforts by many investigators to search for BCDs with equally low or lower metallicities there remained a significant gap in metallicity between I~Zw~18 and all other known BCDs.  This gap was partially filled when Izotov et al.\\ (1990), using the Russian 6--m telescope, tentatively determined the BCD SBS~0335--052 to have a metallicity lower than that of I Zw 18. However, later Melnick et al. (1992) concluded that $Z$ = 1/41 $Z_{\\odot}$ in SBS 0335--052, making it the second most metal--deficient galaxy known in the Universe.  This BCD is now considered as one of the best candidates for a galaxy in formation.   A dwarf galaxy experiencing its first burst of star formation should have an extremely low metallicity in its HII regions. Such a galaxy should not show any sign of an older stellar population produced in previous SF episodes. The gas mass would be expected to significantly exceed that in stars.  Studies of SBS~0335--052 over the last few years show that it appears to fulfill all those conditions and that it is very probably a truly young galaxy. We briefly summarize its main observational properties :  1. The oxygen abundance in the brightest HII region in the galaxy is found to be 12 + log(O/H) = 7.30 $\\pm$ 0.01 (with variations from 7.15 to 7.33 on small scales), only slightly higher than that in I~Zw~18 ( Melnick et al. 1992; Izotov et al. 1997; Izotov et al. 1999).  2. Blue underlying, extended low--intensity emission is detected in SBS~0335--052 on $V$, $R$ and $I$--band images.  The blue $(V-I)$ and $(R-I)$ color distributions suggest that a significant contribution to the emission of the extended low--intensity envelope is due to ionized gas.  It is found that the observed equivalent width of H$\\beta$ emission in the extended envelope is 2 -- 3 times weaker than the value expected in the case of pure gaseous emission.  This could be explained by the presence of underlying stellar emission from coeval A stars (Izotov et al.\\ 1997; Thuan et al.\\ 1997).   3. From the optical colors of the low surface--brightness component together with evolutionary synthesis models, Papaderos et al. (1998) conclude that the age of the underlying stellar population is less than $\\sim$ 100 Myr, with a total mass of $\\sim$ 10$^7$\\,$M_\\odot$, comparable with the total mass of young blue stars (Izotov et al. 1997). Near-infrared (NIR) colors after correction for the gas contribution are consistent with a stellar population not older than 4 Myr and the possible contribution to the NIR light from an evolved stellar population in SBS 0335--052 cannot exceed $\\sim$ 15\\% (Vanzi et al. 2000).  4. No evidence for strong, narrow Ly$\\alpha$ emission is found in a UV spectrum obtained with the {\\sl Hubble Space Telescope} ({\\sl HST}) (Thuan \\& Izotov 1997).  Instead, a damped Ly$\\alpha$ absorption line is observed which implies an HI column density $N$(H\\,I) = 7 $\\times$ 10$^{21}$ cm$^{-2}$, the largest observed so far for a BCD and comparable to the highest column densities observed in Ly$\\alpha$ clouds in the direction of quasars. The resonant O\\,I $\\lambda$1302 line is also detected. Assuming that the O\\,I line is unsaturated and that it originates in the neutral gas leads to an oxygen abundance in the neutral gas of $\\sim$ 1/37000 of the solar value, or about 1/900 of that in the HII regions around the young super--star clusters. However, a more consistent interpretation of the UV absorption lines of O, Si and S is that these lines originate not in the HI but in the HII gas. If this interpretation is correct, then the neutral gas in this system can be pristine and not polluted with heavy elements at all (Thuan \\& Izotov 1997).  5. Thuan et al. (1995) and Izotov \\& Thuan (1999) have also argued on the basis of the very small dispersion of the C/O and N/O ratios in extremely metal--deficient BCDs with a metallicity less than 1/20 solar, that in those galaxies, C and N are made exclusively by massive ($M$ $>$ 9 $M_\\odot$) stars, because intermediate-mass (3 $M_\\odot$ $\\leq$ $M$ $\\leq$ 9 $M_\\odot$) stars have not had time to evolve and release their nucleosynthesis products. Izotov \\& Thuan (1999) suggest that BCDs with $Z$ $<$ $Z_\\odot$/20 are young, with ages less than $\\sim$ 100 Myr, the main--sequence lifetime of a 9 $M_\\odot$ star being $\\sim$ 40 Myr. Chemical abundances thus also suggest that SBS~0335--052, with $Z$ = $Z_\\odot$/40, might be a young galaxy.   In view of the unusual properties of this candidate young galaxy we considered it important to measure its neutral gas content and study the HI structure and velocity field.  Measurements of the integrated HI flux of this galaxy in the 21--cm line were carried out with the Nan\\c cay radio telescope (NRT) in 1993 (Thuan et al.\\ 1999a). From that data the total HI mass was estimated which indicates that this galaxy is gas rich ($M_{\\rm HI}$ $\\sim$ 10$^{9}$ $M_{\\odot}$) with less than a tenth of its baryonic mass in stars.  The HI flux appeared to be high enough to allow higher resolution mapping with the VLA. In this paper we present the results of VLA observations of SBS~0335--052 and discuss the properties of its HI density distribution and velocity field.  The paper is organized as follows. In section 2 we describe the observations and data reduction. In section 3 we derive the HI density distribution and velocity field. In section 4 we examine the gas and dynamical mass of the system and in section 5 we focus on various aspects related to the star formation in SBS~0335--052. The paper is summarized in Section 6.    \\begin{deluxetable}{lcc} \\tablenum{1} \\tablecolumns{3} \\tablewidth{0pt} \\tablecaption{Observational parameters \\label{tbl-1}} \\tablehead{ \\colhead{Parameter} & \\colhead{Value} & \\colhead{Value} } \\startdata Date\t\t\t& 26 Dec 94         \t&   13 May 95\t\t\\nl VLA configuration\t\t&    C              \t&      \t\\nl No.\\ of antennae\t&    26             \t&      26\t\\nl Time on source [min]    &   163                 &      80        \\nl Central velocity [km\\,s$^{-1}$]\t&  4043  \t&    4043\t\\nl Half--power beam\t& 20\\farcs5$\\times$15\\arcsec\\  &  61\\arcsec$\\times$45\\arcsec\\\t\\nl Brightness temperature [K]& 2.0                   &   0.22       \\nl corresponding to 1 mJy\\,beam$^{-1}$\t& \t&\t\t\\nl \\nl Total bandwidth [MHz]\t&    3.12 \t\t&     3.12\t\\nl Total no.\\ channels\t&   31 \t\t\t&    31 \t\\nl Channel width [km\\,s$^{-1}$]\t&   21.2 \t&    21.2 \t\\nl Noise in channel maps [mJy\\,beam$^{-1}$]\t& 0.5\t&     1.0 \\nl \\nl Total bandwidth [MHz]  \t&    1.56\t\t&     1.56\t\\nl Total no.\\ channels \t&   63\t\t\t&    63\t\t\\nl Channel width [km\\,s$^{-1}$]    &    5.3 \t&     5.3\t\\nl Noise in channel maps [mJy\\,beam$^{-1}$] & 1.0\t&     2.0\t\\nl \\enddata \\end{deluxetable} ",
        "conclusions": "  \\begin{enumerate}  \\item{The optical dwarf galaxy SBS~0335--052 is associated with a huge HI cloud with overall size of 66 $\\times$  22\\,kpc, elongated in the EW direction. The cloud has a total HI mass of $1.68 \\times 10^{9}$ $M_{\\odot}$.}  \\item{This HI cloud contains two prominent HI peaks, located nearly symmetrically to the East and West, and separated by about 22 kpc. They measure 7 $\\times$ 4 kpc and contain 0.79 $\\times 10^{9}$ $M_{\\odot}$ and 0.89 $\\times 10^{9}$ $M_{\\odot}$, respectively, of HI.}  \\item{The two HI density peaks are identified with two optical emission--line dwarf galaxies, the eastern one with SBS~0335--052 itself, and the western one with a new, one order of magnitude less luminous dwarf named SBS~0335--052W. This latter galaxy shows very similar properties, including an extremely low metallicity of the HII region gas, consistent with the amount produced during a first SF burst.  Radial velocities of these galaxies, as measured from the emission lines of their HII regions are very close to those of the HI gas at the positions of their respective peaks.}  \\item{The velocity field of the HI cloud is rather disturbed, with presumably tidal tails protruding from both edges of the cloud. This suggests that what we are seeing is an interaction, possibly with the nearby galaxy NGC 1376. This large spiral galaxy, the dominant member of the galaxy group LGG 103, with a projected radial velocity of 120\\,km\\,s$^{-1}$ larger than that of SBS~0335--052, is located at a projected distance of 150 kpc. Consideration of the relevant parameters of the SBS~0335--052 system and NGC~1376 does not contradict the hypothesis that the formation of the two young dwarfs from the huge low--surface density HI cloud was triggered by a tidal interaction with NGC~1376.}  \\item{ Alternatively, we are dealing with two almost equal mass HI clouds, each measuring $\\sim 7 \\times 4 $\\,kpc, each containing an optical counterpart, and showing solid--body rotation with an amplitude of $\\sim$ 35-40 \\,km\\,s$^{-1}$. In this case, the triggering of the star formation in the two HI clouds is probably due to their mutual tidal interaction.   }  \\item{Each component seems to require the presence of dark matter, and the system as a whole is dominated by a dark matter halo. The total dynamical mass of the system is estimated to be between 3.8 $\\times$ 10$^{9}$ $M_{\\odot}$ and 9 $\\times$ 10$^{9}$ $M_{\\odot}$, or $\\sim$ 1.9 to  4.5 times larger than the total mass of gas and stars.}  \\item{If the viewing angle of the HI disk is less than 77$^{\\circ}$, the gas surface densities for both HI peaks are larger than the critical value determined from the parameters of the rotation curve at those radii and are consistent with the onset of gravitational instability and subsequent star formation.}  \\end{enumerate}   {\\it Acknowledgements.} We thank the partial financial support of NATO collaborative research grant 921285 (SAP, TXT), INTAS grant No. 94--2285 (SAP, YII),  NSF grant AST-9616863 (TXT, YII), Russian grant RFBR 97$-$02$-$16755 (SAP) and  CONACyT grant No.\\ 0460P--E (EB). SAP is grateful to the NRAO staff at Socorro for hospitality and support with the VLA data reduction, and to Oleg Verkhodanov and Pat Murphy for help with installing {\\sc aips} at the Special Astrophysical Observatory. The $B$ optical image of the SBS~0335--052 system was kindly provided by Polychronis Papaderos. The anonymous referee made useful suggestions which helped to improve the presentation of the paper. We have made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We have also used the Digitized Sky Survey which was produced at the Space Telescope Science Institute under US Government grant NAG W-2166."
    },
    "0011/astro-ph0011014_arXiv.txt": {
        "abstract": "We present a new phenomenological model for the spectral energy distribution of normal star-forming galaxies between 3 and 1100 $\\mu$m.  A sequence of realistic galaxy spectra are constructed from a family of dust emission curves assuming a power law distribution of dust mass over a wide range of interstellar radiation fields.  For each interstellar radiation field heating intensity we combine emission curves for large and very small grains and aromatic feature carriers.  The model is constrained by \\IRAS\\ and ISOCAM broadband photometric and ISOPHOT spectrophotometric observations for our sample of 69 normal galaxies; the model reproduces well the empirical spectra and infrared color trends.  These model spectra allow us to determine the infrared energy budget for normal galaxies, and in particular to translate far-infrared fluxes into total (bolometric) infrared fluxes.  The 20 to 42 \\m\\ range appears to show the most significant growth in relative terms as the activity level increases, suggesting that the 20--42 \\m\\ continuum may be the best dust emission tracer of current star formation in galaxies.  The redshift dependence of infrared color-color diagrams and the far-infrared to radio correlation for galaxies are also explored. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011222_arXiv.txt": {
        "abstract": "Astronomy is entering a new era as multiple, large area, digital sky surveys are in production. The resulting datasets are truly remarkable in their own right; however, a revolutionary step arises in the aggregation of complimentary multi-wavelength surveys (\\ie the cross-identification of a billion sources). The federation of these large datasets is already underway, and is producing a major paradigm shift as Astronomy has suddenly become an immensely data-rich field. This new paradigm will enable \\emph{quantitatively and qualitatively new science}, from statistical studies of our Galaxy and the large-scale structure in the universe, to discoveries of rare, unusual, or even completely new types of astronomical objects and phenomena. Federating and then exploring these large datasets, however, is an extremely challenging task. The Digital Sky project was initiated with this task in mind and is working to develop the techniques and technologies necessary to solve the problems inherent in federating these large databases, as well as the mining of the resultant aggregate data. ",
        "introduction": "The Digital Sky project is an NPACI (National Partnership for Advanced Computing Infrastructure, an NSF Computer Science initiative) sponsored program to study the role of advanced computational systems (both processor and network oriented) in the distribution of data from multiple large area, digital sky surveys. This project was originally conceived by Tom Prince, the principal investigator for the project, in 1996 as way to leverage high performance computing to tackle some of the incipient problems involved in federating and mining the large amounts of Astronomical information that were beginning to become available. The project was initially funded in 1997, and has since grown to include a large number of participants from Caltech Astronomy, the Caltech Center for Advanced Computing Research, the Infrared Processing and Analysis Center, the Jet Propulsion Laboratory, and the San Diego Center for Supercomputing.  One of the fundamental tenets of the Digital Sky project is the requirement not to develop another analysis tool (which is the last thing needed by the Astronomical community). Instead, we conceived that the project would be able to optimally achieve its goals by serving as a technology demonstrator. Initially, we focused on identifying the set of requirements to create a prototype virtual observatory (\\ie the ``Digital Sky''). Afterwards, we researched the core problem of astronomical data federation. Currently, we are exploring the concept of image data mining.   \\begin{figure}[!htb] \\plotfiddle{brunnerr3-1.eps}{3.0 in}{0}{40}{40}{-144}{0} \\caption{ \\small A top-down overview of the architecture for a virtual observatory. The ability to view either a seamless view of the sky or the detailed capabilities of any individual archive is the overriding requirement.} \\end{figure}  Overall, our approach is designed to employ a minimalist design, that focuses more on integrating existing components in cooperation with community experts (see Figure 1 for an overview). Any knowledge or applications developed in the course of this project have been disseminated to the appropriate experts within the community who are in interested in identifying and attempting to solve the complex problems that arise in federating disparate, large area, digital sky surveys. ",
        "conclusions": ""
    },
    "0011/gr-qc0011115_arXiv.txt": {
        "abstract": "s{This talk is based on my work in collaboration with B.~Boisseau, D.~Polarski, and A.A.~Starobinsky. The most natural and best-motivated alternatives to general relativity are the so-called ``scalar-tensor'' theories, in which the gravitational interaction is mediated not only by a (spin-2) graviton, but also by a (spin-0) scalar field. We study quintessence in this general framework, and show that the microscopic Lagrangian of the theory can be unambiguously reconstructed from two observable cosmological functions of the redshift: the luminosity distance and the linear density perturbation of dustlike matter. We also analyze the constraints imposed on the theory by the knowledge of only the first of these functions, as it will probably be available sooner with a good accuracy.}  Theoretical attempts at quantizing gravity or at unifying it with other interactions generically predict the existence of scalar partners to the graviton. For instance, in superstring theory, a dilaton is already present in the supermultiplet of the 10-dimensional graviton, and many other scalar fields, called the ``moduli'', appear when performing a Kaluza-Klein dimensional reduction down to our usual spacetime. Gravity theories incorporating such scalar fields have a great importance in cosmology, notably in inflationary scenarios, but also for ``quintessence'' models, in which the cosmological constant suggested by observations of type-Ia supernovae appears in fact as the present value of a scalar-field potential $U(\\Phi)$.  Starobinsky\\,\\cite{St98} proved, in the case of a minimally-coupled scalar field, that the precise shape of this potential can be reconstructed if one knows either the present energy density of dustlike matter $\\Omega_{m,0}$ and the luminosity distance $D_L(z)$ as a function of the redshift $z$, or the Hubble constant $H_0$ and the density perturbation of dustlike matter $\\delta_m(z)$ as a function of $z$. However, the models inspired by extra-dimensional theories generically involve a non-minimal coupling between the scalar field and the curvature. Several parametrizations of the action may be chosen,\\cite{beps00,ep00} but it can always be written in the following ``Brans-Dicke'' form: \\begin{equation} S = \\int d^4 x \\sqrt{-g}\\left\\{\\Phi R - {\\omega(\\Phi)\\over \\Phi}\\left(\\partial_\\mu\\Phi\\right)^2-2 U(\\Phi)\\right\\} + S_{\\rm matter}\\left[{\\rm matter}; g_{\\mu\\nu}\\right]\\ , \\end{equation} where ``matter'' denotes all kinds of matter fields, including gauge bosons (we consider here theories satisfying the ``weak equivalence principle''). Since this action involves two unknown functions of the scalar field, $\\omega(\\Phi)$ and $U(\\Phi)$, we now need both observable functions $D_L(z)$ and $\\delta_m(z)$ to fully reconstruct the theory. On the other hand, the present values of $\\Omega_{m,0}$ and $H_0$ are not necessary.  The first step consists in obtaining the Hubble function $H(z)$ from the knowledge of $D_L(z)$, like in general relativity since this is purely kinematical: $1/H(z) = [D_L(z)/(1+z)]'$, where a prime denotes a derivative with respect to $z$ (and where a corrective factor involving $\\Omega_{\\kappa,0}\\equiv -\\kappa/(a_0H_0)^2$ enters if the universe is spatially curved\\,\\cite{ep00}). The perturbation equations in the limit of short wavelengths and the background Friedmann-Robertson-Walker equations then give \\begin{eqnarray} {\\Phi\\over \\Phi_0} &\\simeq& {3\\over2}\\left({H_0\\over H}\\right)^2 {(1+z)\\,\\Omega_{m,0}\\,\\delta_m\\over \\delta_m''+[H'/H - (1+z)^{-1}]\\delta_m'} \\times \\left(1+{1\\over 2\\omega+3}\\right)\\ , \\label{eq2}\\\\ {2U\\over (1+z)^2H^2} &=& \\Phi'' + \\left({H'\\over H}- {4\\over 1+z}\\right)\\Phi' +\\left[{6\\over (1+z)^2} - {2\\over 1+z}\\,{H'\\over H} - 4 \\left({H_0\\over H}\\right)^2\\! \\Omega_{\\kappa,0}\\right]\\Phi \\nonumber\\\\ && - 3\\,(1+z)\\left({H_0\\over H}\\right)^2\\Phi_0\\,\\Omega_{m,0}\\ , \\label{eq3}\\\\ \\omega &=& -{\\Phi\\over \\Phi'^2} \\Biggl\\{\\Phi'' + \\left({H'\\over H} + {2\\over 1+z}\\right)\\Phi' -2\\left[{1\\over 1+z}\\,{H'\\over H} - \\left({H_0\\over H}\\right)^2\\! \\Omega_{\\kappa,0}\\right]\\Phi \\nonumber\\\\ && + 3\\,(1+z) \\left({H_0\\over H}\\right)^2 \\Phi_0\\,\\Omega_{m,0} \\Biggr\\}\\ , \\label{eq4} \\end{eqnarray} where the integration constant $\\Phi_0$ may be set to 1 without loss of generality. The three functions $\\Phi(z)$, $U(z)$ and $\\omega(z)$ --and thereby $U(\\Phi)$ and $\\omega(\\Phi)$-- can thus be reconstructed unambiguously if $D_L(z)$ and $\\delta_m(z)$ are experimentally determined precisely enough. Since solar-system experiments constrain the present value of $\\omega$ to be larger than 2500, and that it can be proven\\,\\cite{beps00} to be larger that $\\sim 5$ even for redshifts $z\\sim1$, one can further simplify Eq.~(\\ref{eq2}) by suppressing its last factor inside parentheses. Note that in that case, the three equations become algebraic.  As $D_L(z)$ will be observed sooner with a good accuracy, we analyzed the constraints imposed on scalar-tensor theories by this single function.\\cite{ep00} We assumed particular forms for either $\\omega(\\Phi)$ or $U(\\Phi)$, and reconstructed the other one thanks to Eqs.~(\\ref{eq3})-(\\ref{eq4}) above. Our main conclusion is that the knowledge of $D_L(z)$ over a wide redshift interval, say up to $z \\sim 2$, is sufficient to distinguish these theories from general relativity plus a cosmological constant, even if there are large (tens of percents) experimental errors. The reason for this strong result is that we took into account not only solar-system (and binary-pulsar) constraints, but also the following important {\\it theoretical\\/} constraints: $\\bullet$~The graviton should carry positive energy ($\\Rightarrow \\Phi > 0$). $\\bullet$~The scalar field should carry positive energy ($\\Rightarrow 2\\omega(\\Phi)+3 > 0$). $\\bullet$~The potential $U(\\Phi)/\\Phi^2$ should be bounded by below and have a reasonable shape, and to get a stable theory, the scalar mass squared should be positive ($\\Rightarrow \\{\\Phi[2\\omega(\\Phi)+3]^{-1/2}[U(\\Phi)/\\Phi^2]'\\}'\\geq 0$, where a prime denotes here $d/d\\Phi$). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011152_arXiv.txt": {
        "abstract": "We present long-slit spectroscopy, B and R bandpass imaging, and 21 cm observations of a sample of early-type galaxies in nearby clusters which are known to be either in a star-forming phase or to have had star formation which recently terminated.  From the long-slit spectra, obtained with the Blanco 4-m telescope, we find that emission lines in the star-forming cluster galaxies are significantly more centrally concentrated than in a sample of field galaxies.  The broadband imaging reveals that two currently star-forming early-type galaxies in the Pegasus I cluster have blue nuclei, again indicating that recent star formation has been concentrated.  In contrast, the two galaxies for which star formation has already ended show no central color gradient. The Pegasus I galaxy with the most evident signs of ongoing star formation (NGC7648), exhibits signatures of a tidal encounter.  Neutral hydrogen observations of that galaxy with the Arecibo radiotelescope reveal the presence of $\\sim$4 x $10^8 M_{\\sun}$ of HI.  Arecibo observations of other current or recent star-forming early-type galaxies in Pegasus I indicate smaller amounts of gas in one of them, and only upper limits in others.  The observations presented above indicate that NGC7648 in the Pegasus I cluster owes its present star formation episode to some form of tidal interaction.  The same may be true for the other galaxies with centralized star formation, but we cannot rule out the possibility that their outer disks have been removed via ram pressure stripping, followed by rapid quenching of star formation in the central region. ",
        "introduction": "It is now well-established that the star formation rate (SFR) in rich clusters of galaxies has declined markedly since z$\\sim$0.5.  Evidence for this decline was first established by Butcher \\& Oemler (1978, 1984; see Margoniner \\& de Carvalho 2000 and references therein for recent results), who found a systematic increase in the fraction of blue galaxies in rich clusters with increasing redshift, a phenomenon known as the Butcher-Oemler effect.  {\\it HST} images have since shown that the increase in the population of blue galaxies with z is also accompanied by an increase in the fraction of spiral and interacting/merging galaxies.  Thus the decline in SFR in rich clusters since z=0.5 appears connected to the decrease in the population of spirals, and consequent increase in S0's (e.g., Dressler \\etal 1997; Poggianti \\etal 1999), though the details are complex (Fabricant \\etal 2000).  The aim of this paper is to better understand what are the driving mechanisms that produce the observed falloff in SFR.  While the basic observation of the Butcher-Oemler effect appears to be generally well-accepted, there is much current debate concerning the causes of the declining SFR, and to what extent the rich cluster environment is involved.  On the one hand, the cluster environment is certainly capable of both provoking and quenching star formation in the member galaxies, since both gas removal mechanisms (e.g., Gunn \\& Gott 1972; Abadi, Moore, \\& Bower 1999) and tidal perturbations (e.g., Lavery \\& Henry 1988; Byrd \\& Valtonen 1990; Moore \\etal 1996, 1998; Bekki 1999), which can lead to inward flow of gas and subsequent enhanced star formation, have been proposed.  Evidence for enhanced star formation comes from the discovery of galaxies in distant clusters with strong Balmer absorption lines and no emission lines in their spectra (e.g., Dressler \\& Gunn 1983). This so-called ``E + A'' or ``K + A'' phenomenon is often attributed to the effect of a recent starburst (e.g., Barger et al. 1996), thus indicating that cluster galaxies are ending their star-forming careers in a final burst.  Recently, early-type galaxies in {\\it nearby} clusters, such as Coma, with similar K+A spectra have been discovered with surprising frequency by Caldwell \\etal (1993) and Caldwell \\& Rose (1997).  Subsequent long-slit spectroscopy (Caldwell \\etal 1996) and {\\it HST} imaging (Caldwell, Rose, \\& Dendy 1999) have indicated that the recent star formation in these galaxies has been centrally concentrated to the inner $\\sim$2 kpc in radius.  In addition, models of the integrated spectra of some of these systems indicate that the recently terminated star formation was burst-like, not simply a truncation of ongoing star formation (Caldwell \\etal 1996; Caldwell \\& Rose 1998). Thus in these cases the evidence favors a centrally concentrated star formation episode, rather than truncated disk star formation.  On the other hand, the identification of the K+A phenomenon with a starburst episode has been challenged by Newberry, Boroson, \\& Kirshner (1990) and by Balogh \\etal (1999).  These authors have pointed out that the distribution of emission line and Balmer line strengths in cluster galaxies does not demonstrably differ from what one might expect for a collection of fading spirals whose star formation was truncated rather suddenly.  Furthermore, Balogh \\etal (1999) propose that the fraction of K+A galaxies in distant clusters is actually far lower (only about 1\\% of their sample) than previously indicated, and that there is no indication that the fraction of K+A's exceeds that in the field, nor that the fraction has evolved significantly from z=0.5 to the present epoch.  This result, based on the CNOC sample, contrasts strongly with the results of the MORPHS sample, which is reported in Dressler \\etal (1999) and Poggianti \\etal (1999).  One problem is that, except for immediately after a burst episode when the colors are very blue and the Balmer lines are very strongly enhanced, a fading starburst and a fading, truncated disk produce very similar spectra. Thus exceptional care (high signal-to-noise) and fortune are required to distinguish the two. As a result, there is considerable uncertainty as to whether the rather widely assumed poststarburst picture for K+A galaxies is indeed correct, rather than the less dramatic situation of a truncated spiral.  A key unresolved question is the matter of the rapid evolution in SFR since z=0.5: Is it  caused by final large bursts of star formation in cluster galaxies, or is it instead produced by a less dramatic truncation of star formation?   Previous efforts to distinguish between the two scenarios in the first question have largely concentrated on interpreting the star formation histories of galaxies through their colors and spectra.  However, another way to resolve the issue is by investigating the {\\it spatial distribution of the last star formation} in these galaxies.  If the primary effect is simply termination of star formation in an otherwise normal spiral disk, then we would expect to see the fading of a roughly uniformly blue galaxy, with the exception, perhaps, of a redder bulge.  If the mechanism is somehow related to starbursts, however, one would expect the fading blue population to be centrally concentrated, since current hypotheses for the production of starbursts all utilize variations on a theme of tidal gravitational perturbations, which cause a pileup of gas in the centers of galaxies, and hence a central starburst.  In this paper we have combined spectroscopic and imaging data of star forming galaxies in nearby clusters to evaluate whether recent star formation in cluster galaxies has been centrally concentrated. We place particular emphasis on the perspective of nearby clusters, since while in distant clusters the processes producing the rapid evolution in SFR are more in evidence, in nearby clusters one can study the processes with higher spatial resolution and higher S/N ratio, both in spectroscopy and imaging.  Two datasets relevant to unusual star formation in nearby clusters are presented in this paper.  First, we investigate the spatial distribution of ionized gas in 10 early-type galaxies in nearby rich clusters.  We compare the central concentration of the ionized gas with that in an extensive sample of low-redshift field galaxies. The purpose is to evaluate whether or not star formation in the cluster galaxies is unusually centrally concentrated when compared with the field sample.  Second, we present B and R images and HI observations of four early-type galaxies in the Pegasus I cluster for which Vigroux \\etal (1989) have previously found evidence for ongoing or recently completed star formation.  The Pegasus I cluster is particularly relevant here because of its low cluster x-ray emission, which indicates that ram pressure stripping is not a significant process in this cluster.  Thus we have an unusual opportunity to isolate the effects of external (tidal) perturbations on the star formation histories in cluster galaxies.  The basic plan for this paper is as follows. In \\S 2 the imaging and spectroscopic data are summarized.  In \\S 3 we compare the spatial distribution of ionized gas in early-type galaxies in nearby rich clusters with that of the sample of ``field'' galaxies.  In \\S 4 we discuss radial luminosity profiles, morphology, and $B-R$ radial color maps of four early-type galaxies in the nearby Pegasus I cluster, as well as 21 cm observations of these galaxies. In \\S 5 we combine the results of the previous sections to assess the nature of the last star formation in cluster galaxies, and place our results in the context of other recent analyses. ",
        "conclusions": "To summarize, two principal results have been obtained.  First, long-slit spectroscopy of emission-line galaxies in the Coma, DC2048-52, and DC0326-53/DC0329-52 nearby rich clusters has demonstrated that centralized (in comparison to a field galaxy sample) episodes of star formation occur in galaxies inhabiting the cluster environment at the present epoch.  Second, we have found, on the basis of B and R imaging (viz., the presence of reverse color gradients in their central regions), that centralized star formation is also occurring in galaxies in the Pegasus I cluster.  The key point regarding the Pegasus I cluster is that it represents a mild cluster environment in which ram pressure stripping is unlikely to be a factor.  We now elaborate on these two results, placing them into context with other research on the evolution of galaxies in clusters.  \\subsection{Starbursts versus Truncated Star Formation}  A number of recent investigations have addressed the issue of whether star formation in cluster galaxies primarily ends through rapid truncation or through a major starburst, a question that has been debated for some time (e.g., Newberry, Boroson, \\& Kirshner 1990).  Specifically, studies of the color-magnitude relations of galaxies in relatively distant clusters (e.g., Kodama \\& Bower 2000), as well as the statistics of emission line equivalent widths (e.g., Balogh \\etal 1997), indicate that many, if not most, of the galaxies are losing their gas through a truncation of disk star formation, rather than via a major outburst of star formation. However, the evidence presented here, along with the evidence of centralized star formation in cluster galaxies already published in Caldwell \\etal (1996) and Caldwell, Rose, \\& Dendy (1999), demonstrates that in at least some cluster galaxies, centralized episodes of star formation have taken place which are very different from ``normal'' star formation in a disk.  Thus some mechanism must be sought that either drives gas to the central regions of galaxies and provokes star formation there, or that selectively removes gas from the outside and then produces a terminal episode in the center.  As is discussed in \\S 5.2 below, in the Pegasus I cluster the evidence points quite clearly in the direction of a tidally induced star formation episode rather than gas removal by ram pressure.  Thus the cluster environment is capable of producing both truncated star formation and enhanced centralized outbreaks as well.  It is natural to suspect, therefore, that both evolutionary mechanisms are important in clusters.  To place the evidence for centralized star formation found here within the context of other work, we mention two other results which may be related. Koopmann \\& Kenney (1998) have emphasized the peculiar aspects of the morphologies of spirals in the Virgo cluster.  In particular, they discuss a class of ``truncated'' spirals (which they refer to as morphological type St), that have nuclear emission and central light concentrations typical of late-type spirals, but quiescent disks that are usually associated with early-type spirals. Hence the Virgo St spirals, which have centrally concentrated emission, may be analogs to the early-type galaxies with emission that we are finding in other, more distant, clusters. In addition, Moss \\& Whittle (1993, 2000) have found large numbers of spiral galaxies in 8 nearby clusters which have ``circumnuclear'' (i.e., centrally concentrated) emission.  Moreover, the spirals with circumnuclear emission are associated with morphologies indicative of tidal interaction. Thus it appears likely that our cluster emission-line galaxies have much in common with those of Moss \\& Whittle.  \\subsection{Tidal Disturbances versus Ram Pressure Stripping}  Our second result, obtained from the Pegasus I cluster, addresses the role of gas stripping versus tidal interaction in driving galaxy evolution in clusters.  Various lines of evidence (cf., discussion in \\S 1) have convincingly demonstrated the importance of ram pressure stripping in depleting the gas reservoirs in galaxies in the rich cluster environment.  The primary observational evidence for ram pressure stripping {\\it in action} comes from studies such as Gavazzi \\etal (1995) and Kenney \\& Koopmann (1999), where the radio and optical morphologies show that star formation is preferentially occurring in an asymmetric extranuclear arc.  The inference is that star formation is induced at the interface between the leading edge of the infalling galaxy and the hot cluster ICM.  The observational signatures have now been modelled by Quilis, Moore, \\& Bower (2000), whose 3-d gas dynamical simulations also show the characteristic bow shock effect. On the other hand, in the mild environment of the Pegasus I cluster, ram pressure stripping cannot plausibly be a factor, and yet two early-type galaxies are found with ongoing centralized star formation.  In addition, the morphology of the most clearcut star forming case in Pegasus I, NGC7648, is different from those of the three galaxies in A1367 which are believed to be in dynamical interaction with the cluster ICM.  While the morphology of NGC7648 is complex, there is no sign of a preferential edge (bow shock) to the star formation.  Rather, non-spiral disturbances are concentrated into knots and perhaps ripples, which are suggestive of tidal interaction.  Thus our second result is that tidal disturbances do indeed play an important role in triggering star formation episodes in the cluster environment.  Evidence for the role of tidal disturbances in clusters has been presented in several previous studies.  Rubin, Waterman, \\& Kenney (1999) have demonstrated that approximately half of the spiral galaxies in the Virgo cluster have disturbed rotation curves.  While they argue that the disturbances are likely to be tidal in nature, they do not distinguish between galaxy-galaxy and galaxy-(sub)cluster interactions.  A specific example of an early-type Virgo spiral believed to be a late-stage merger remnant is NGC4424 (Kenney \\etal 1996). Moss and Whittle (2000) find that a high percentage of the circumnuclear Emission line galaxies in nearby clusters have perturbed morphologies indicative of destabilization through a tidal disturbance of some kind. Numerous studies of more distant clusters also provide strong evidence for the importance of tidal interactions and mergers.  Ground-based imaging by Thompson (1988), Lavery \\& Henry (1988), and Lavery, Pierce, \\& McClure (1992) revealed numerous cases of interactions in distant clusters, a result that has been dramatically confirmed by {\\it HST} imaging (e.g., Dressler 1994a,b; Couch \\etal 1998 and references therein).  \\subsection{Triggering Mechanism for Centralized Star Formation}  Finally, a key question, which remains to be addressed, is what kind of tidal (or perhaps other) disturbance triggers the centralized star formation episodes which are clearly evident in the cluster galaxies studied in this paper.  Due to its presence in a ram-pressure-free environment, NGC7648 offers a particularly good perspective on this question. However, the comparison of NGC7648 with four late-stage merger candidates in \\S 4.4 leads to inconclusive results.  With its low (and likely centrally concentrated) HI content, and apparent lack of an exponential disk component, NGC7648 fails to fit in with either the equal mass or intermediate mass ratio merger scenarios.  The one late-stage merger candidate with similarly low HI content, NGC4424, is resident in the Virgo cluster, and believed to be a victim of ram pressure stripping (Kenney \\etal 1996).  It appears that either the progenitor galaxies (perhaps low in atomic hydrogen) and/or encounter geometry must be fundamentally different from the late-stage mergers considered here, or the merger scenario might be incorrect.  In that case, another type of tidal disturbance, e.g., ``galaxy harassment'' (Moore \\etal 1996; 1998) or destabilization from the rapidly varying (sub)cluster gravitational potential (Byrd \\& Valtonen 1990; Bekki 1999), could supply the triggering mechanism.  Unfortunately, the data presented here is not sufficient to resolve this important issue.  In conclusion, we have shown that some early-type galaxies in nearby clusters are experiencing, or have recently experienced, unusual star formation histories when compared with those in more isolated galaxies. Specifically, star formation in these galaxies is more centrally concentrated than in the disks of isolated spiral galaxies.  Thus the cluster environment is clearly influencing the evolution of these galaxies. The cause of this centralized star formation is somewhat less clear.  In principle it could be due to the preferential removal, by ram-pressure stripping, of the atomic gas in the galaxy's disk, thereby allowing only for star formation in the central region.  Alternatively, it may result from external tidal perturbations, which cause gas to be driven into the central region of the galaxy.  The fact that in some cases there is good evidence that the star formation rate has been substantially enhanced in a recent starburst favors the tidal triggering mechanism."
    },
    "0011/astro-ph0011478_arXiv.txt": {
        "abstract": "Based upon multi-epoch \\it Hubble Space Telescope \\rm observations, we present the discovery of sixteen high-quality Cepheid candidates in NGC~4527.  Corrected for metallicity effects in the Cepheid period-luminosity relation, we derive a distance, including both random (r) and systematic (s) uncertainties, of 13.0$\\pm$0.5(r)$\\pm$1.2(s)\\,Mpc. Our result is then used to provide a calibration of the peak $B$-, $V$-, and $I$-band luminosities of the peculiar Type~Ia supernova SN~1991T, a resident of NGC~4527.  Despite its documented spectroscopic peculiarities, after correction for the decline rate--luminosity correlation the corrected peak luminosity is indistinguishable from those of so-called ``normal'' Type~Ia SNe.  With now nine local calibrators at our disposal, we determine a robust value for the Hubble Constant of H$_0$=73$\\pm$2(r)$\\pm$7(s)\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "introduction": "\\label{introduction}  When corrected for the shape of their light curves, the peak apparent brightnesses of Type~Ia supernovae (SNe) provide an outstanding secondary distance indicator (\\eg, Phillips et~al. 1999; Jha et~al. 1999; Saha et~al. 1999; Gibson et~al. 2000a).  First, the extreme luminosity of the SNe allows one to probe substantially further into the unperturbed Hubble flow, in comparison with competing secondary indicators such as the Tully-Fisher relation and surface brightness fluctuations.  Second, the intrinsic scatter is smaller than that of other indicators.  While attempts have been made to provide a purely theoretical calibration for the peak luminosity (\\eg, H\\\"oflich \\& Khokhlov 1996), remaining model uncertainties still favor the observational approach.  The \\it Sandage/Tammann/Saha Type~Ia SN HST Calibration Project \\rm (Saha et~al. 1999, and references therein) was designed to provide just such a calibration, by determining the Cepheid distances to eight nearby galaxies (IC~4182, NGC~5253, 4536, 4496A, 4639, 3627, 3982, 4527) host to Type~Ia SNe.  This original program has been supplemented with two additional calibrators situated within NGC~3368 (Tanvir et~al. 1999) and NGC~4414 (Turner et~al. 1998).  Table~1 of Gibson et~al. (2000a) provides a complete list of the galaxies and SNe covered by all these calibration programs, including a subjective quality ranking.  Of the ten nearby Type~Ia SNe targeted by the various teams, only one is classified as spectroscopically peculiar --- SN~1991T.  While the data suggest that SN~1991T was a supernova with peculiar surface abundances (no evidence for Si, Ca, or S absorption during the early phases of the explosion), and had a light curve in $B$ and $V$ that was marginally broader than usual, in most other ways it seems to have been a normal Type~Ia explosion (Phillips et~al. 1992).  Derivation of the Cepheid distance to NGC~4527, the host galaxy to SN~1991T, provides a direct test of the robustness of the technique as applied to potentially peculiar SNe.  In what follows, we present the Cepheid distance to NGC~4527, provide a calibration for the peak $B$-, $V$-, and $I$-band luminosities of SN~1991T, and present an improved determination of the Hubble Constant. As will be seen, the derived peak luminosities do not appear to be affected by the spectroscopic peculiarities of the SN. ",
        "conclusions": "\\label{summary}  A Cepheid-based distance to NGC~4527, host to the peculiar Type~Ia SN~1991T, has been derived using the same software pipeline and unbiased Cepheid PL analysis employed throughout the HST Key Project on the Extragalactic Distance Scale series of papers. The corrected peak luminosity is indistinguishable from that of spectroscopically normal SNe, demonstrating the robustness of the corrected peak luminosity as a secondary distance indicator.  SN~1991T is only the fourth calibrator SN for which accurate $I$-band photometry exists.  NGC~4527, in conjunction with re-derived distances for eight other Type~Ia SN-host galaxies, and a full accounting of random and systematic uncertainties, yields a robust value of the Hubble Constant of H$_0$=73$\\pm$2(r)$\\pm$7(s)\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.   \\newpage"
    },
    "0011/astro-ph0011364_arXiv.txt": {
        "abstract": "The possibility of exciting the g-modes of a solar-type star as a consequence of the gravitational interaction with a close companion (a planet or a brown dwarf) is studied by a perturbative approach. The amplitude of the emitted gravitational wave is computed and compared with the quadrupole emission of the system, showing that in some cases it can be considerably larger. The effects of radiation reaction are considered to evaluate the timescale of the emission process, and a Roche lobe analysis is used to establish the region where the companion can orbit without being disrupted by tidal interactions with the star. ",
        "introduction": "In recent years  a significant number of systems composed by a solar-type star and one or more orbiting companions has been discovered in our neighbourhood. According to the {\\it Extrasolar Planets Catalog}\\cite{Uff2}, at the time we write this paper  58 extrasolar planetary systems have been confirmed, and this number is destinated to grow in the near future (see \\cite{review} for a recent review). 46 of the observed orbiting objects are planets with masses ranging between [0.16-11]$M_J$, where $M_J$ is Jupiter's mass; 12 are bigger, up to  $60~M_J$, and include  super-planets and brown dwarfs. In the following, we shall indicate all the orbiting objects as `planets'.  The newly discovered systems exhibit some unexpected features; 22 planets  move on elliptic orbits with eccentricities $e > 0.3$, larger than the largest  eccentricity encountered in  the solar system (Mercury, $e = 0.2$), and  many of them are incredibly close to the central star, much closer than expected. For instance the planet orbiting  51 Pegasi, discovered in 1995\\cite{51Pega},\\cite{51Pegb}, has a period $P=4.23$ days and orbits its star at a distance  $a=0.05$ AU. More than two thirds of the planets discovered up to now are orbiting their host star much closer than Mercury orbits the Sun ($P=88$ days,   $a=0.39$ AU), and even closer than 51 Pegasi's planet. These orbital properties  challenge the existing theories on planet formation and evolution, and suggest an interesting possibility: there may exist planets on  orbital radii sufficiently small to excite some g-modes of the star by tidal interaction. This issue has been discussed by Terquem et al. \\cite{terquem} in connection with the short orbital period of 51 Pegasi's planet. They calculated the dynamical tides raised by the planet on the star, and the dissipation timescales due to turbulent viscosity in the convective zone and to radiative damping in the radiative core. In this paper we shall approach the problem of the excitation of the g-modes from a different point of view: by using a perturbative approach, justified by the fact that the mass of the central star is much bigger than that of the planet, we shall compute the gravitational signal emitted by a  system in which the planet  moves on a circular orbit of radius ${R}_0$, close to that which would correspond to the resonant excitation of a g-mode,  $R_{res}$. We shall extend the work presented in a previous paper (\\cite{bdf}, to be referred to hereafter as Paper I), were we have studied the quadrupole emission due to the orbital motion of the  observed planetary systems, and started to  explore the  possibility for a planet to excite  some stellar modes. First of all, in section II we shall verify by a Roche-lobe analysis whether the planet can move on the orbit  $R=R_0$ without being disrupted by tidal interactions. In section  III,  using the quadrupole formalism we shall compute the characteristic amplitude of the gravitational signal, $h_{Q}$, emitted by the system because of the orbital motion  of the planet  on  ${R}_0$. As usual, the two bodies will be treated as pointlike masses, and the effects of the planet on the internal structure of the star will be neglected. These effects will be considered in section IV, where  we shall assume that the planet induces a small perturbation on the gravitational field of the star and on its thermodynamical structure. We shall  integrate the perturbed Einstein's equations coupled to  the hydrodynamical equations, having the stress energy tensor of the planet as a source,  and compute the characteristic amplitude of the emitted gravitational signal, $h_{R}$, which will be compared with $h_{Q}$. In this way we will be able to evaluate the enhancement in the radiation emitted by the system as a consequence of the g-modes excitation. In Section V we shall discuss how long can a planet  move on an orbit close to resonance before radiation reaction effects lead the planet off resonance. Conclusions will be drawn in section VI. ",
        "conclusions": "The study carried out in this paper has been motivated by the recent discovery of several extrasolar planetary systems very close to Earth, with companions orbiting the host star much closer than previously expected. This discovery suggests the possibility that a companion may be found orbiting at such close distance  to excite one of the g-modes of the star, and a Roche-lobe analysis indicates that this may be possible, especially if the companion is a brown dwarf. The frequency of a g-mode in a sun-like star is of the order of a few digits in $10^{-4}$ Hz, and would be in the bandwith of the spaceborn interferometer LISA. Thus, our aim was to understand how much such resonant process may enhance the emission of gravitational waves of the system with respect to the radiation it emits because of the time varying quadrupole moment associated to the orbital motion. This enhancement has been  evaluated by integrating the equations that describe the perturbation the companion induces on the star, and  computing the amplitude of the emitted wave when the planet moves close to an orbit corresponding to the resonant excitation of a g-mode. The amplitude increases sharply as the orbit approaches the resonant one; however, in order to get a chance to observe one of these signals emitted by a system in our vicinity, it is also needed that  it stays above a certain threshold long enough. For this reason we have also computed the effects of radiation reaction on the shrinking of the orbit and the associated timescales. The main information we extract from  our  study  is that only close brown dwarf companions may produce signals of some relevance;  a brown dwarf of  40 Jovian masses, orbiting a star located at a distance of 10 pc from Earth, if close to an orbit resonant with the mode $g_4$, may emit radiation of amplitude greater than $\\sim 2\\cdot 10^{-20}$ at a frequency of the order of $\\sim  10^{-4}$ Hz,  spanning a region $\\Delta R=1.3$ km over a time interval of $\\sim 4$ years. Of course the amplitude of the wave increases with the mass of the brown dwarf  and decreases with the distance,  and it will be interesting to see whether astronomical observations will identify any  such system in our vicinity in the future.  As a last remark, we would like to mention that things would not be significantly different had we considered more realistic models of solar-type stars, such as those discussed in \\cite{christensen},\\cite{dirty}; they predict oscillation frequencies that are  a little higher than those of our simple polytropic model; however, our frequencies are  at least qualitatively correct to give order-of-magnitude estimates for the effects of resonant excitation of the stellar modes and  for the excitation timescales."
    },
    "0011/astro-ph0011370_arXiv.txt": {
        "abstract": "We present analysis of WFPC2 imaging of two spiral galaxies partially backlit by E/S0 systems in the pairs AM1316-241 and AM0500-620, as well as the (probably spiral) foreground system in NGC 1275. Images in $B$ and $I$ are used to determine the reddening curve of dust in these systems. The foreground spiral component of AM1316-241 shows dust strongly concentrated in discrete arms, with a reddening law very close to the Milky Way mean ($R=E_{B-V}/A_V= 3.4 \\pm 0.2$). The dust distribution is scale-free between about 100 pc and the arm dimension, about 8 kpc. The foreground spiral in AM0500-620 shows dust concentrated in arms and interarm spurs, with measurable interarm extinction as well. In this case, although the dust properties are less well-determined than in AM1316-241, we find evidence for a steeper extinction law than the Milky Way mean (formally, $R \\approx 2.5 \\pm 0.4$, with substantial variation depending on data quality in each region). The shape of the reddening law suggests that, at least in AM1316-241, we have resolved most of the dust structure. In AM0500-620 it is less clear that we have resolved most of the dust structure, since the errors are larger. In AM0500-620, the slope of the perimeter-scale relation (associated with fractal analysis) steepens systematically on going from regions of low to high extinction. A perimeter-smoothing length test for scale-free (fractal) behavior in AM1316-241 shows a logarithmic slope typically $-0.4$ on 100-1000 pc scales. However, we cannot determine a unique fractal dimension from the defining area-perimeter relation, so the projected dust distribution is best defined as fractal-like. For scales above 2-4 pixels (120-250 pc), the box-counting estimate gives a fractal dimension close to 1.4, but the perimeter-area relation gives a dimension of 0.7 on large scales and inconsistent results for small scales, so that the distribution shows only some aspects of a fractal nature. In neither galaxy do we see significant regions even on single-pixel scales in spiral arms with $A_B > 2.5$. The measurements in NGC 1275 are compromised by our lack of independent knowledge of the foreground system's light distribution, but masked sampling of the absorption suggests an effective reddening curve much flatter than the Milky Way mean (but this may indicate that the foreground system has been affected by immersion in the hot intracluster gas, or is inside the stellar distribution of NGC 1275). The bright blue star clusters trace the absorption in this system quite closely, indicating that these clusters belong to the foreground system and not to NGC 1275 itself. ",
        "introduction": "The importance of dust extinction within galaxies is a fundamental problem for a wide range of issues. It affects the extragalactic distance scale through inclination corrections to the Tully-Fisher relation. Dust extinction affects the inferred evolution of quasars (the high-redshift cutoff, in particular) via cumulative extinction by foreground galaxies along typical lines of sight. Estimates of star-formation rates and mass-to-light ratios in individual galaxies are strongly dependent on estimates of internal absorption in galaxies. Assessing the ``typical\" extinction within spiral galaxies has proven controversial, and various approaches give substantially different results (and the various interpretations of these results diverge even more widely). Although the presence of obscuring matter was clear from the earliest systematic galaxy photography (\\cite{curtis18}), only much later was the connection between inclination and surface brightness used to estimate how much internal extinction there might be (\\cite{holmberg58}). Holmberg's reassuring conclusion, that internal extinction is not a major factor in the global emerging light from spiral galaxies, eventually came to be challenged on two grounds. Disney et al. (1989) showed that the blue colors and inclination behavior of spiral disks could be consistent with models of high optical depth, so that what we see is dominated by a small proportion of unreddened starlight. Valentijn (1990) drew attention to this issue by analyzing the surface brightness and inclination of spirals in the ESO-LV catalog (\\cite{esolv}), and claiming that these quantities are empirically so independent that typical spirals must be optically thick to blue light across much of their visible disks. This analysis (though not necessarily its conclusion) was challenged by Burstein et al. (1991), who attributed Valentijn's main result to possible selection effects in the galaxies included in the ESO-LV sample, and asserted that the problem could only be solved when truly volume-limited galaxy samples could be assembled.  Because of the assumptions required in using these model-based techniques to estimate extinctions, we have pursued a more direct approach: observing overlapping galaxy pairs to measure dust extinction in foreground spirals via differential surface photometry. As set out in White \\& Keel (1992), the ideal case for such an analysis has a foreground spiral half projected against a similarly sized background elliptical. This geometry allows us to recover estimates of the light from each galaxy by itself, using symmetry from the non-overlapping portions of the images, and thence to reconstruct the absorption in the overlapping part of the foreground spiral, limited only by the level of asymmetry in the spiral (since elliptical and S0 systems are so precisely symmetric). An extensive ground-based program (\\cite{wkc}, hereafter WKC; \\cite{wkc96}; \\cite{kw}) identified the best candidates for such analysis, and measured disk opacities for 10 spiral galaxies. The general results were:  -- spiral arms and resonance rings can have large optical depths ($\\tau_B > 2$) over a wide range of galactocentric distance.  -- interarm dust has a roughly exponential distribution with radius, and causes only mild extinction in the outer disk (typically $A_B \\geq 0.5$ only within 0.4$R_{25}$).  -- the extinction ratios from $B$ to $I$ passbands are greyer than the mean Milky Way curve, which we interpreted as evidence of patchy or clumpy extinction on unresolved size scales.  This last point made it especially desirable to measure dust properties on finer spatial scales, perhaps recovering the intrinsic reddening curve of the material and facilitating a comparison between grain populations in various spiral galaxies. This was pursued using two approaches. We describe here high spatial resolution observations achieved by HST imaging of two especially amenable galaxy pairs, AM1316-241 and AM0500-620 (plus archival data on the backlit high-velocity system in NGC 1275). A complementary analysis, comparing far-IR and submillimeter emission with optical absorption measures for three overlapping galaxy pairs (including AM1316-241) was presented by Domingue et al. (1999).  These galaxy pairs were selected as the most promising cases, for symmetry and angular size, from the initial survey reported by WKC. The two newly-observed pairs were both initially catalogued by Arp \\& Madore (1987). AM1316-241 (ESO 508-IG45) consists of a highly inclined spiral in front of an elliptical galaxy with very nearly circular isophotes, as measured from the nonoverlapping regions. The spiral has measured redshift $cz=10365$ km s$^{-1}$, notably larger than that of the elliptical (9700 km s$^{-1}$, White \\& Keel 1992; Donzelli \\& Pastoriza 2000). AM0500-620 (ESO 119-IG27) consists of a nearly face-on spiral of type Sb partially projected in front of a smooth system with a light profile suggesting an S0 classification. Donzelli \\& Pastoriza list heliocentric velocities of 9005 and 8811 km s$^{-1}$ for the spiral and elliptical components, respectively. For linear scales, we adopt a Hubble constant of 75 km s$^{-1}$ Mpc$^{-1}$, for which a single WFPC2 $0.1\\arcsec$ pixel subtends 65 and 57 pc, respectively, in AM1316-241 and AM0500-620. We also consider archival WFPC2 data on the object in the foreground of NGC 1275; in this case the foreground system has a redshift $cz=8200$ km s$^{-1}$, while NGC 1275 itself has $cz=5264$ km s$^{-1}$. Given their context in the Perseus/Abell 426 cluster, the line-of-sight separation of these two galaxies is almost indeterminate. Some suggest that the two galaxies are strongly interacting, perhaps enhancing the luminosity of the emission-line filaments around NGC 1275 proper (\\cite{hu83}), and absorption geometry indicates that the foreground system  must lie at least outside these filaments (\\cite{keel83}). For the PC CCD, which encompasses most of the foreground absorption, the image scale in this case is 15 pc/pixel, although we worked with mosaiced data at 30 pc/pixel. ",
        "conclusions": "We have used HST imaging of two spiral galaxies partially backlit by elliptical and S0 systems to provide direct measures of extinction and dust structure of the foreground spirals. This approach provides extinction estimates which are independent of the assumptions about internal structure that are necessary to perform many of the statistically-based opacity tests often applied to this question.  Opacity maps show an interesting contrast between the spiral components of AM1316-241 and AM0500-620. In both cases, the arm/interarm dichotomy in extinction is strong. However, the statistics of coverage at various transmission levels are notably different in the two galaxies. The arm/interarm dichotomy in extinction is strengthened at the high resolution afforded by HST; much of what appeared in an earlier ground-based study as diffuse interarm dust is now seen to be comprised of narrow spurs crossing the main spiral pattern (as in AM0500-620). For the areas we can analyze, located at radii of $0.4-0.8 ~ R_{25}$, the dust is heavily concentrated to arms, with half the dust mass contained in only 20\\% of the projected area and 95--98\\% of the dust mass contained in half the area. Kiloparsec-scale complexes of absorbing material are estimated to contain $3-7 \\times 10^4$ $M_{\\sun}$ of dust grains, with correspondingly large implied gas masses, $1-2 \\times 10^7$ $M_{\\sun}$, if we assume a typical Galactic gas-to-dust ratio. Interarm extinction is quite small, although the statistical distribution by area differs for the two spirals. There is a well-defined low-extinction peak in the areal distribution for AM1316-241, while the distribution of extinction values is much broader and continuous with the arm regions in AM0500-620.  Reddening curves have been measured by comparing transmitted light pixel by pixel between $B$ and $I$ passbands. For AM1316-241, we find a well-determined slope given by $R = A_V/E_{B-V} = 3.4 \\pm 0.2$, close to the Galactic value of 3.1 derived largely from diffuse regions. The data quality for AM0500-620 is poorer, and the overall fit to a Galactic law (or any plausible reddening law) is poor. For the single cloud in AM0500-620 with the best data, we do see reasonable reddening behavior, with a formal fit of $R=2.5 \\pm 0.4$. In NGC 1275, our accuracy is limited by our inability to estimate the brightness of the foreground galaxy. We have used a color criterion to eliminate pixels with obvious foreground light and find the remainder have a reddening curve much flatter (greyer) than the local norm. This might imply a different dust composition, a different fine-scale distribution, or that the foreground system is so close to NGC 1275 that it is at least partially embedded within its stellar distribution (so that our simple geometric scheme for measuring opacity breaks down).  We speculate, moreover, that the dust grain size distribution in the NGC 1275 foreground system has been truncated at the low end by sputtering in the hot intracluster medium of Abell 426. Lack of particles smaller than $\\sim 0.5 \\mu$ would flatten the extinction curve effectively. Following Draine \\& Salpeter (1979), the survival time scale for particles of radius $a$ embedded in a hot medium of density $n$ is $$t_{sp}  =  10^6~{a \\over n} ~{\\rm years},$$ for $a$ in $\\mu$m and $n$ in cm$^{-3}$. For a representative value of $n = 10^{-3}$, the timescale for $a = 0.25 \\mu$ is $10^9$ years. At the high velocity difference between this system and NGC 1275 proper (itself essentially at rest in the cluster) it would have moved at least 3 Mpc during the time it would take to sputter a Galactic grain population to this extent. This is comparable to the extent of X-ray emission from the Perseus cluster (\\cite{ettori}). A strict application of this sputtering prescription would weight the distribution toward larger grains only for very specific size distributions, since for spherical grains and a power-law size distribution, the net effect of sputtering would be to shrink all grains at a constant rate by area (rather than volume). In this case, a given set of original grains would move down the size distribution in step, with the only net affect being progressive loss of the largest grains. However, recent modeling of sputtering effects on more realistic grain structures suggest that the net effect can indeed by a greying of the extinction curve (\\cite{aguirre99a}). Furthermore, preferential ejection of small grains due to either ram pressure or galactic winds would weight the distribution toward larger sizes (\\cite{aguirre99b}). If the grain population in the foreground system of NGC 1275 has indeed been altered, the original dust mass and typical $B$ extinction would have been substantially larger than we observe (with the $I$ extinction being less affected, since larger grains are more important at this wavelength). If the system has interacted with the intracluster medium, the small H~I mass (\\cite{vge83}) is easier to reconcile with the morphological evidence for a late Hubble type, since H~I is the first ISM component lost from cluster spirals (\\cite{cayatte}; \\cite{bravo}).  We have used the perimeter-smoothing length test to see whether the dust in these galaxies can be well described by fractal (scale-free) structures, as has been found for H~I distributions in nearby spirals. We find scale-free behavior for the perimeter of a given contour in dust column density as a function of smoothing scale from 0.05--8 kpc. However, we cannot determine a fractal dimension from these data, because the scaling between different contours (at different transmission levels) at a given smoothing length does not show the requisite linear relation for log (area) versus log (perimeter) relation (the perimeter in fact peaks within the measured area range). We can thus describe the dust distribution as fractal-like but not formally fractal, with box-counting dimension close to 1.4 but area-perimeter measures giving a dimensions near 0.7. The box-counting dimension is within the range found for H I structures in galaxies by Westpfahl et al. (1999), but given the discrepancy between the two dimensions and the non-monotonic area-perimeter relation, the dust distribution cannot be accurarely described as fractal. The area covered at a given extinction level is nearly independent of smoothing scale over the range we can measure, so that our dust mass estimates are not biased by structures smaller than our resolution.  What processes define and control fine structure in the dust? Grains are generally strongly coupled to other components of the interstellar medium, both empirically and for robust theoretical reasons, so the dust distribution will reflect a range of physical processes that need not impact the grains directly. Strictly dynamical features of disk dynamics can concentrate the ISM into regions of high density, and can change the configuration of these regions on dynamical timescales (a recently discussed example is amplification of small fluctuations as seen in the nuclear dust patterns of NGC 2207; \\cite{elm98}). Energy input from stellar winds and supernovae has a crucial role in altering the small-scale structure of the ISM (Rosen \\& Bregman 1995), both mechanically and by driving phase changes (including grain evaporation in the most extreme environments). These data will not support a detailed analysis of such issues, but the differences in dust distributions may most plausibly be linked to changes in the energy-input rates and perhaps to different dominant mechanisms shaping the ISM in these two galaxies."
    },
    "0011/astro-ph0011146_arXiv.txt": {
        "abstract": "We present $V-$band surface photometry and major-axis kinematics of stars and ionized gas of three early-type spiral galaxies, namely NGC 772, NGC 3898 and NGC 7782.  For each galaxy we present a self-consistent Jeans model for the stellar kinematics, adopting the light distribution of bulge and disc derived by means of a two-dimensional parametric photometric decomposition.  This allowed us to investigate the presence of non-circular gas motions, and derive the mass distribution of luminous and dark matter in these objects.  NGC 772 and NGC 7782 have apparently normal kinematics with the ionized gas tracing the gravitational equilibrium circular speed. This is not true in the innermost region ($|r|\\la8''$) of NGC 3898 where the ionized gas is rotating more slowly than the circular velocity predicted by dynamical modelling.  This phenomenon is common in the bulge-dominated galaxies for which dynamical modelling enables us to make the direct comparison between the gas velocity and the circular speed, and it poses questions about the reliability of galaxy mass distributions derived by the direct decomposition of the observed ionized-gas rotation curve into the contributions of luminous and dark matter. ",
        "introduction": "\\label{sec:introduction}  In the past kinematic and dynamical studies of disc galaxies have been mainly focused on late-type spirals. With respect to lenticular and early-type spiral galaxies (defined as those with $B/T>0.3$, following Simien \\& de Vaucouleurs 1986), the study of late-type spirals has many advantages from both the observational and interpretative points of view. Late-type spirals have small or even negligible bulges, so they are usually described as luminous discs embedded in dark matter halos.  They are on average more gas-rich than early-type spirals. Gas is mainly confined in the galactic disc, considered moving at circular speed and therefore used as direct tracer of the gravitational potential. Since the line-of-sight velocity distribution (LOSVD hereafter) of ionized or neutral hydrogen is easier to measure than the stellar one, the number of spirals of known gaseous kinematics increased faster than that of the galaxies for which stellar motions were measured.  These facts made late-type spirals easy targets, particularly in the search for dark matter (see Sofue \\& Rubin 2001 for a review).  Only recently Heraudeau \\& Simien (1998) and Heraudeau et al. (1999) measured the stellar rotation curves and velocity-dispersion profiles along the major axis of a large sample of spirals (more than 60 objects) ranging from Sa to Sc.  In last years systematic surveys have been started to derive the stellar and gaseous kinematics in S0's (Bertola et al. 1995; Fisher 1997; Bettoni \\& Galletta 1997) and in early-to-intermediate type spirals (Corsini et al. 1999; Vega Beltr\\'an 1999; Vega Beltr\\'an et al. 2000).  If the direct comparison of stellar and gaseous kinematics allowed to unveil also in disc galaxies the presence of kinematically-decoupled components (see Bertola \\& Corsini 1999 for a recent review), the application of dynamical models showed that in the bulge-dominated region the ionized-gas velocity may fall below the circular speed (Fillmore, Boroson \\& Dressler 1986; Kent 1988; Kormendy \\& Westpfahl 1989). These `slowly-rising' rotation curves have been interpreted as the signature of the presence of pressure-supported gas (Bertola et al. 1995; Cinzano et al. 1999), which has been recently expelled by stars of the bulge but not yet heated to the virial temperature of the galaxy. This discovery poses new questions about the reliability of galaxy mass distributions derived by the decomposition of the observed ionized-gas rotation curve into the contributions of luminous and dark matter. In fact the inner gradient of the gas velocity curve is usually used to fix the amount of luminous matter, which results to be underestimated if gas velocity rises slowly than circular speed.  In this paper we present a study of the ionized-gas and stellar kinematics for 3 Sa -- Sb galaxies, namely NGC 772, NGC 3898, and NGC 7782. We apply a self-consistent dynamical model based on a two-dimensional photometric decomposition method to explain the different kinematical behaviour of gas and stars.  The goal of this work is to derive the mass distribution of these objects and investigate the possible presence of non-circular gas motions in the bulge region.  The paper is organized as it follows. In Sect.~\\ref{sec:properties} we give an overview of the properties of NGC 772, NGC 3898 and NGC 7782; in Sect.~\\ref{sec:observations} we present the broad and narrow-band imaging and major-axis kinematics of stars and ionized gas; in Sect.~\\ref{sec:structure} we discuss the photometric and dynamical techniques adopted to investigate the light and mass distribution of these galaxies; in Sect.~\\ref{sec:results} we apply these techniques to the galaxies, deriving their structural and kinematic parameters; finally, Sect.~\\ref{sec:conclusions} is devoted to the discussion of the results and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Non-circular velocities of the ionized gas in the bulge-dominated region of early-to-intermediate type disc galaxies have been detected by different authors (Fillmore et al. 1986; Kent 1988; Kormendy \\& Westpfahl 1989). The discovery that in a large fraction of S0's the gas velocity dispersion remains as high as the stellar velocity dispersion over an extended radial range (Bertola et al. 1995) corroborated the early suggestions that at small radii gas could be dynamically hot rather than settled in a cold disc. In this framework, the dynamical modelling by Cinzano et al. (1999) showed that in the S0 galaxy NGC 4036 the difference between the circular velocity curve inferred from stellar kinematics and the rotation curve of the gas is substantially accounted for its high velocity dispersion.  It is usually customary, in measuring mass distribution, to assume that the ionized gas is moving in circular orbits. It is therefore crucial to understand if and when the warm gas in galaxies has a rotation curve which rises more slowly than the gravitational equilibrium circular velocity. Ignoring this effect can lead to a significant underestimate of the masses of the bulge, disc and dark halo (e.g. Kent 1988 for a discussion).  Up to now such a direct comparison between the circular velocity inferred from dynamical modelling based on stellar kinematics and surface photometry, and the observed gas rotation velocity has been done for only a small set of lenticulars and early-to-intermediate spirals (see Table \\ref{tab:slowlyrising}). The three new cases discussed in this paper represent therefore a useful step in understanding how common the phenomenon of slowly-rising rotation curves is in bulge-dominated galaxies.  For NGC 772, NGC 3898 and NGC 7782 we present a self-consistent Jeans model for the stellar kinematics, adopting the light distribution of bulge and disc derived by means of a two-dimensional parametric photometric decomposition of the observed $V-$band surface brightness.  This allows us not only to investigate the presence of non-circular gas motions, but also to derive the mass distribution of luminous and dark matter in these objects.  \\begin{table*} \\caption{Early-to-intermediate type disc galaxies with available comparison between gas rotation velocity and circular speed} \\begin{center} \\begin{tabular}{lllrccrccccl} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{Object} & \\multicolumn{2}{c}{Morphological Type} & \\multicolumn{1}{c}{$D$} & \\multicolumn{1}{c}{$M^0_{B_T}$}  & \\multicolumn{1}{c}{$M_{\\it bulge}/M_{\\it disc}$}  & \\multicolumn{1}{c}{$r_e$} & \\multicolumn{1}{c}{$r_\\mu$} & \\multicolumn{1}{c}{$V_{\\it gas}<V_{\\it circ}$} & \\multicolumn{1}{c}{$r_V$} & \\multicolumn{1}{c}{$r_\\sigma$} & \\multicolumn{1}{c}{Ref.} \\\\ \\noalign{\\smallskip} \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[RSA]} & \\multicolumn{1}{c}{[RC3]} & \\multicolumn{1}{c}{[Mpc]} & \\multicolumn{1}{c}{[mag]}& \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[kpc]} & \\multicolumn{1}{c}{[kpc]} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[kpc]} & \\multicolumn{1}{c}{[kpc]} & \\multicolumn{1}{c}{} \\\\ \\noalign{\\smallskip} \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & \\multicolumn{1}{c}{(8)} & \\multicolumn{1}{c}{(9)} & \\multicolumn{1}{c}{(10)} & \\multicolumn{1}{c}{(11)} & \\multicolumn{1}{c}{(12)} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC~772  & Sb(rs)I      & .SAS3..   & 34.7 & -21.61 & 2.11 & 11.0 & 6.2 & no  & -      & 0.5 &1, $V$\\\\ NGC~2179 & Sa           & .SAS0..   & 33.7 & -19.81 & 2.80 &  2.0 & 5.7 & no  & -      & 0.2 &2, $R$\\\\ NGC~2775 & Sa(r)        & .SAR2..   & 13.9 & -19.91 & 1.39 &  3.6 & 4.4 & yes & 2.0    & -   &2, $r$\\\\ NGC~3898 & SaI          & .SAS2..   & 17.1 & -19.84 & 3.55 &  1.6 & 2.3 & yes & 0.7    & -   &1, $V$\\\\ NGC~4036 & S0$_3$(8)/Sa & .L..$-$.. & 20.1 & -20.03 & 2.04 &  1.2 & 0.7 & yes & 1.0    & 0.4 &3, $V$\\\\ NGC~4450 & Sab pec      & .SAS2.    & 17.0 & -20.40 & 2.69 &  1.7 & 1.4 & yes & $>5.0$ & ?   &4, $B$\\\\ NGC~4569 & Sab(s)I-II   & .SXT2..   & 17.0 & -21.36 & 0.08 &  0.3 & 0.4 & yes & 2.6    & ?   &4, $B$\\\\ NGC~5055 & Sbc(s)II-III & .SAT4..   &  7.5 & -20.35 & 0.12 &  0.4 & 0.3 & yes & 0.5    & ?   &4, $B$\\\\ NGC~7782 & Sb(s)I-II    & .SAS3..   & 74.8 & -21.95 & 0.24 &  1.1 & 1.5 & no  & -      & 0.4 &1, $V$\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} \\label{tab:slowlyrising} \\end{tabular} \\begin{minipage}{16.cm} NOTE -- col. 2: morphological type from RSA; col. 3: morphological type from RC3; col. 4: distance derived as $V_0/H_0$ with $V_0$ the velocity relative to the centroid of the Local Group obtained from the heliocentric systemic velocity as in RSA and $H_0 = 75$ \\kms\\ Mpc$^{-1}$. The heliocentric velocity of NGC 5055 has been taken from RC3. For NGC 4450 and NGC 4569 which belong to the Virgo Cluster the distance has been taken following Freedman et al. (1994); col. 5: absolute corrected $B$ magnitude from $B^0_T$ in RC3; col. 6: bulge-to-disc mass ratio; col. 7: effective radius of the $r^{1/4}$ bulge. For NGC 7782 the bulge has an exponential surface-brightness profile; col. 8: radius at which $\\mu_{\\it bulge} = \\mu_{\\it disc}$; col. 9: yes = the ionized-gas component shows a slowly-rising rotation curve (i.e. $V_{\\it gas} < V_{\\it circ}$ at small radii), no = the gas rotation traces circular speed at all radii ($V_{\\it gas} = V_{\\it circ}$ at all radii); col. 10: radial range in which $V_{\\it gas} < V_{\\it circ}$; col. 11: radial range in which $\\sigma_{\\it gas} \\simeq \\sigma_{\\it stars}$, ? = the radial profile of the ionized-gas velocity dispersion is not available; col. 12: references and bands of the data [1 = this paper, 2 = Corsini et al. (1999), 3 = Cinzano et al. (1999), 4 = Fillmore et al. (1986)]. \\end{minipage} \\end{center} \\end{table*}   In NGC~772 the observed velocity dispersion of the ionized gas is comparable to the stellar velocity dispersion and is far higher than expected from the thermal motions or small-scale turbulence, typically observed in a gaseous disc. However, dynamical modelling showed that the rotation of the ionized-gas traces the circular velocity and the central rise observed in its velocity dispersion is due to seeing effects, which serves as a warning against over interpretation.  On the contrary, the gas rotation curve of NGC 3898 can be classified as `slowly-rising', since for the inner $8''$ ionized gas is rotating more slowly than the circular velocity predicted by dynamical modelling.  NGC 3898 is the only galaxy of our sample for which previous determinations of the masses of the bulge and disc components are available in the literature (Kent 1988; Fillmore et al. 1986; Moriondo et al. 1998b). Comparison shows that NGC 3898 is a clearcut example that mass decomposition based only on emission-line rotation curves can be unreliable, at least for bulge-dominated galaxies. This error was made and discussed by Kent (1988) and Moriondo et al. (1998b), who assumed the gas to be rotating at circular velocity at all radii and found extraordinarily low values for the mass-to-light ratio of the spheroidal component, as if NGC 3898 was an almost `bulge-less' galaxy in spite of its overall morphology (see Table \\ref{tab:properties} and Figs. \\ref{fig:n3898_kinematics}, and \\ref{fig:n3898_Halpha}) and bulge-to-disc luminosity ratio ($B/D=2$).  The asymmetry of the gas rotation curve of NGC 3898 (with a maximum $\\Delta V$ between the two sides of about 80 \\kms\\ at $6''$ from the centre), and the strong difference between the gas ($\\sigma_{\\it gas} \\leq 90$ \\kms) and the stellar velocity dispersion ($\\sigma_{\\it stars} > 120$ \\kms\\ with a central peak to $210$ \\kms) in the innermost $5''$ (where we observe almost the same velocity gradient for both gas and stars) are an indication that along with random motions other phenomena could also be contributing to the slowly rise of the gas velocity. This is the case of other galaxies with slowly-rising rotation curves, such as the Sa NGC 2775 (Corsini et al. 1999) and the S0 NGC 4036 (Cinzano et al. 1999). The asymmetry of the inner parts of the gas rotation curve of NGC 2775 has been explained as due to a gaseous component which is not rotating in the galaxy plane, and the possible presence of drag forces between the ionized gas and the hot component of the interstellar medium of NGC 4036 has been suggested (for a discussion of the phenomenon see Cinzano et al. 1999).  Finally, in NGC~7782 the gas rotation curve shows small deviations from pure circular rotation only in the transition region between bulge and disc.  Any evidence for possible non-circular motions is marginal and the sharp increase of the gas velocity dispersion towards the nucleus can be better explained as due to a seeing effect as in the case of NGC 772.  As far as the presence of dark matter is concerned, we infer that the mass is essentially traced by light in NGC 772 and NGC 7782, where gas rotation velocities were observed out to $0.2\\,R_{25}$ and $0.6\\,R_{25}$, respectively. For NGC 3898 the combined stellar and gaseous rotation data (which extend out to $0.2\\,R_{25}$ and $0.7\\,R_{25}$ respectively) require the presence of a massive dark halo. This result is based on the idea that gas kinematics at large radii is representative of a dynamically cold disc supported by rotation, as suggested by the high gas rotation velocity and low velocity dispersion ($V/\\sigma\\simeq10$) for $r>30''$ (corresponding to about $0.2\\,R_{25}$). This result also qualitatively agrees with the general dark matter scenario, with the less massive galaxies being the more dark-matter dominated (Salucci \\& Persic 1999), although these data do not permit a careful measurement of the dark halo properties.  Kent (1988) found that, out of the 14 Sa galaxies he modeled, six show gas rotation curves that rise too slowly for a constant mass-to-light ratio implying they are respectively too low for the bulges and too high for the disc components. According to Table~\\ref{tab:slowlyrising}, the gas motion is not circular in the inner regions of 6 of the 9 early-to-intermediate disc galaxies, for which dynamical modelling allows the direct comparison between the gaseous and circular speeds (among these only NGC 3898 belongs also to Kent's sample).   Without pretending to draw any statistical conclusion based on so small number of objects, it nevertheless seems probable that non-circular gas velocities could be a common feature in the central kiloparsec of disc galaxies ranging from S0's to Sb's. However, according to the available data, it is not possible to derive a straightforward correlation with the galaxy morphological type, total luminosity, bulge-to-disc mass ratio or bulge size. It seems that even if the gas may be supported by non-circular random motions, this type of dynamical-pressure support may not be the only effect for the `slowly-rising' gas rotation curves. In fact, the central gas velocity dispersion is too low to account for the difference between gas and circular speeds in two of the three galaxies (namely NGC 2775, and NGC 3898) shown in Table \\ref{tab:slowlyrising} for which the gas velocity falls short of the circular speed and the velocity-dispersion profile of the gaseous component is available. Dynamical modelling based both on stellar and gaseous kinematics of a larger number of bulge-dominated galaxies is needed to understand the links between slowly-rotating gas and bulge properties, as a first step to put constraints on its physical nature (e.g. Mathews 1990) and origin (e.g. Bertola et al. 1995).  \\medskip \\noindent {\\bf Acknowledgements.}  \\noindent We are grateful to Prof. P.A. Strittmatter, Director of the Steward Observatory and to Dr. G.V. Coyne, S.J., Director of the Vatican Observatory for the allocation of time for our observations. WWZ acknowledges the support of the grant 7914 of the {\\sl Jubil\\\"aumsfonds der Oesterreichischen Nationalbank}. This work was partially supported by grant PB97-0214 of the Spanish DGES."
    },
    "0011/astro-ph0011236_arXiv.txt": {
        "abstract": "There is growing observational evidence that gamma-ray bursts (GRBs) are powered by black holes accreting rapidly through a disk.  The supernova-like outburst that accompanies some gamma-ray bursts suggest that {\\it some} long-duration GRBs may be driven by the accretion of a rotating stellar core onto a central black hole. Such a system can be produced when a compact remnant spirals into the core of its binary companion.  During the inspiral, orbital angular momentum is injected into the helium core.  By the time the compact remnant reaches the center of the helium core, it too has gained angular momentum as well as mass, producing a rapidly accreting black hole (or neutron star) at the center of a rotating stellar core.  Whether or not such a merger (termed He-merger) can produce a GRB depends upon the initial mass and spin of the central black hole, as well as the angular momentum in the stellar core.  In this paper, we use a 3-dimensional smooth particle hydrodynamics (SPH) code to follow the He-merger process and make quantitative estimates of the initial mass and spin of the central compact remnant, as well as the angular momentum in the accreting helium core.  During the inspiral, a $2\\,\\Msun$ compact remnant gains $\\sim 0.5-3.5\\,\\Msun$ depending on the mass of the helium core (more massive cores provide more accretion). The accretion rates on the central remnant are initially very high, $10^5\\,\\Msyr$ up to $10^6\\,\\Msyr$ (the accretion rate increases with increasing helium star mass), and the central remnant quickly becomes a black hole if it was not one already.  From these accretion rates, we estimate GRB explosion energies. In all mergers, magnetically driven jets are expected to  produce GRB explosions with energies above $10^{51}$\\,ergs. For neutrino-annihilation-driven explosions, the GRB energy increases dramatically with helium star mass: the merger of a $2\\,\\Msun$ compact remnant with a $4\\,\\Msun$ helium star only produces a $10^{47}\\,$erg explosion in $\\sim 500$\\,s whereas the merger of a $2\\,\\Msun$ compact remnant with a $16\\,\\Msun$ helium star produces a $>10^{52}\\,$erg explosion in $\\sim 65$\\,s. ",
        "introduction": "\\label{sec:intro}  Although there is no agreed upon engine powering gamma-ray bursts(GRBs), models based upon black hole accretion disks (BHAD models) are attractive.  The progenitors of this class of GRB engine include the mergers of compact binaries (double neutron star, black hole and a neutron star, white dwarf and a black hole), collapsars, and the inspiral of a neutron star or black hole into a helium star (see Fryer, Woosley, \\& Hartmann 1999a for a review). In all of these cases, these progenitors produce accretion disks around black holes which, as they accrete, use neutrinos or magnetic fields to convert part of the gravitational potential or rotational energy of the disk material (or black hole) into a GRB outburst.  The BHAD systems formed by these progenitors differ mainly in the disk mass, disk angular momentum, or amount of material above the plane of the rotation disk and these variations may well explain the diversity of the observed sample of gamma-ray bursts.  For instance, classical GRBs seem to divide into long and short duration subgroups (e.g. Kouveliotou et al. 1993; Katz \\& Canel 1996; Tavani 1998), and it is thought that these two distinct subgroups are explained by different BHAD progenitors (e.g. Fryer et al. 1999a).  Three progenitors may produce long-duration GRBs:  mergers of white dwarfs with black holes, collapsars, and He-mergers. The disks formed in the mergers of white dwarfs with black holes (or neutron stars\\footnote{Belczy\\'{n}ski, Bulik, \\& Rudak (2000) have found that white dwarf/neutron star mergers, which may be able to produce GRB outbursts using a magnetic-field driven mechanism, dominate the number of white dwarf mergers.}) represent one extreme in disk characteristics. Of the three long-duration BHAD GRB progenitors, white dwarf mergers produce disks with the most angular momentum and the least amount of material along the rotation axis.  Although the high angular momentum in these disks make neutrino-mediated explosions extremely weak, white dwarf/black hole mergers may produce strong GRB explosions via some magnetic field mechanism (Fryer et al. 1999b).  Collapsars, on the other hand, represent the other extreme:  they have less angular momentum and much more material along the axis of rotation (MacFadyen \\& Woosley 1999; MacFadyen, Woosley ,\\& Heger 2000). Since the connection between GRB 980425 and supernova 1998bw (Kulkarni et al. 1998), which is naturally explained by collapsar models (Woosley 1993; MacFadyen \\& Woosley 1999; MacFadyen et al. 2000), the collapsar model has become the favorite progenitor for long-duration bursts.  But can all long-duration gamma-ray bursts be collapsars? There is increasing evidence to suggest that collapsars cannot explain all GRBs.  First, the  X-ray and optical lightcurves of GRBs suggest that the environment around the GRB engine is sometimes best fit by a wind-swept model, and other times is best fit by a constant density model (Chevalier \\& Li 2000, Livio \\& Waxman 2000).  Based on the fact that collapsars occur only in massive stars, Chevalier \\& Li (2000) and Livio \\& Waxman (2000) both argue that those GRBs best fit by constant density environments must not be collapsars\\footnote{However, bear in mind that collapsars form only with stars that had very low mass-loss rates (presumably low metallicities) and the winds of collapsar progenitors may not alter the environment significantly (Fryer et al. 1999a). Even those GRBs which are best matched by constant density environments may yet be collapsars.}.  In addition, collapsars may have trouble explaining the iron emission lines detected in three GRBs:  GRB 970508 (Piro et al. 1999), GRB 970828 (Yoshida et al. 1999), and GRB 991216 (Piro, et al. 2000).  To explain these detections (the detections in GRB 970508 and GRB 970828 are marginal at best so not all may be real), B\\\"ottcher (2000) found that the progenitors of these bursts may need to have metallicities in excess of solar. However, the winds of collapsar progenitors must be lower than typical winds of massive stars, and hence collapsar progenitors are likely to be low-metallicity and it is unlikely that collapsars can explain such high iron abundances.  To compare the He-merger model to that of the collapsar, let us review the formation scenario of He-merger GRBs.  The progenitor of this GRB engine is a binary system consisting of two massive stars.  The more massive star (primary) collapses to form a neutron star or black hole.  The loss of matter in the supernova explosion and a possible kick on the compact remnant unbinds many of the binaries.  Those that remain bound tend to be in tight orbits, and when the companion star evolves off the main sequence, it engulfs the compact remnant and the system goes into a common envelope phase.  Once in a common envelope, the compact remnant spirals down through the hydrogen star.  If the binding energy of the envelope is too large, the compact star will spiral into the companion's helium core. It is this inspiral which spins up the helium core and produces the conditions necessary to form a gamma-ray burst.  The merger of the compact remnant with its helium companion provides a mechanism to produce BHAD system very similar to that of a collapsar (Fryer \\& Woosley 1998).  The primary difficulty with the collapsar model is producing a black hole with a rotating stellar envelope.  The helium merger model naturally forms an accreting black hole surrounded by a {\\it rotating} stellar core.  The differences in the formation processes of these two GRB models produce distinguishing features which can be observed.  For example, after the formation of the compact remnant, the He-merger binary system will have some net velocity with respect to its formation region, and it may travel beyond the wind and supernova ejecta of the primary.  Hence, the constant density environment required by some of the lightcurves is easily explained by He-mergers.  In addition, the primary supernova's explosion may enrich the hydrogen envelope of its companion (Israelian et al. 1999). When this envelope is ejected during the common envelope phase, it produces an iron rich torus which may explain the observed iron lines (B\\\"ottcher \\& Fryer 2000).  Fryer \\& Woosley (1998) argued that the engine driving He-merger GRBs would be very similar to collapsars.  But to truly understand the He-merger engine, we must follow the inspiral and obtain quantitative estimates of both the rotation of the helium core and the remnant mass in the core.  Knowing these values, we can then infer the duration and total energy of the He-merger GRB. In this paper, we present hydrodynamical simulations of the inspiral of a compact remnant into a helium core for a range of helium core masses.  We describe our numerical techniques in \\S~\\ref{sec:nm} and the results of our simulations in \\S~\\ref{sec:res}.  We conclude with a discussion of the inspiral results and of the viability of He-mergers as a GRB model. ",
        "conclusions": ""
    },
    "0011/hep-ph0011190_arXiv.txt": {
        "abstract": " ",
        "introduction": "The QED corrections to the scattering process $\\nu_{e} + e \\rightarrow \\nu_{e} + e$ were studied long ago by Lee and Sirlin using the effective four--fermion V$-$A Lagrangian \\cite{LS}, and shortly afterwards Ram \\cite{Ram} extended their calculations by including hard photon emission. A few years later, 't Hooft computed the lowest order prediction to this differential cross section in the Standard Model (SM) \\cite{tH}. Since then, the radiative corrections to this process, which plays a fundamental role in the study of electroweak interactions, have been investigated by many authors, focusing on various aspects of the problem \\cite{ZKN,SU,GV,MS,AHKKM,AH,MSS,DB1,DB2,BBBS,Be2000}.   't Hooft's early SM predictions for $\\nu$--$e$ scattering were used by Bahcall to examine the total cross section, energy spectrum and angular distribution of recoil electrons resulting from the scattering with solar neutrinos \\cite{B87}. More recently, Bahcall, Kamionkowski and Sirlin performed a detailed investigation of the radiative corrections to these recoil electron spectra and total cross sections \\cite{BKS}. Their results show the importance of these corrections for the analysis of precise solar $\\nu$--$e$ scattering experiments, particularly of those measuring the higher energy neutrinos that originate from \\B8 decay.   In this paper we study the $O(\\alpha)$ QED corrections to neutrino electron scattering in the SM, with contributions involving either neutral currents (as in the $\\nu_{\\mu,\\tau} + e \\rightarrow \\nu_{\\mu,\\tau} + e$ process) or a combination of neutral and charged currents (as in the $\\nu_{e} + e \\rightarrow \\nu_{e} + e$ process). In this analysis we make the approximation of neglecting terms of $O(q^2/\\mw^2)$, where $q^2$ is the squared four-momentum transfer and $\\mw$ is the $W$ boson mass.  Within this approximation, which is excellent for present experiments ($|q^2/\\mw^2| \\!\\sim\\! 1$ when the electron recoil energy $\\sim 6 \\times 10^3$ TeV!), the SM radiative corrections to these processes can be naturally divided into two classes. The first, which we will call ``QED'' corrections, consist of the photonic radiative corrections that would occur if the theory were a local four--fermion Fermi theory rather than a gauge theory mediated by vector bosons; the second, which we will refer to as the ``electroweak'' (EW) corrections, will be the remainder. The EW corrections have been studied by several authors \\cite{SU,GV,MS,AHKKM,MSS,BKS} and are not discussed in the present paper. The split--up of the QED corrections is sensible as they form a finite (both infrared and ultraviolet) and gauge--independent subset of diagrams. We refer the reader to ref.~\\cite{Si78-80} for a detailed study of this separation.   The QED radiative corrections are due to both loop diagrams (virtual corrections) and to the bremsstrahlung radiation (real photons) accompanying the scattering process. Of course, only this combination of virtual and real photon corrections is free from infrared divergences. To order $\\alpha$, the bremsstrahlung events correspond to the inelastic process $\\nu_l + e \\rightarrow $ $\\nu_l + e +\\gamma $ ($l=e,\\mu$ or $\\tau$). Experimentally, bremsstrahlung events in which photons are too soft to be detected are counted as contributions to the elastic scattering $\\nu_l + e \\rightarrow$ $\\nu_l + e$. The cross section for these events should be therefore added to the theoretical prediction of the elastic cross section, thus removing its infrared divergence.   We will divide the bremsstrahlung events into ``soft'' (hereafter SB) and ``hard'' (hereafter HB), according to the energy of the photon being respectively lower or higher than some specified threshold $\\epsilon$. We should warn the reader that the words ``soft'' and ``hard'' may be slightly deceiving. Indeed, if $\\epsilon$ is large (small), the SB (HB) cross section will also include events with relatively high (low) energy photons.  While calculations of both soft and hard bremsstrahlung are often performed under the assumption that $\\epsilon$ is a very small parameter, much smaller than the mass of the electron or its final momentum, we will also discuss results for the case in which $\\epsilon$ is an arbitrary parameter constrained only by the kinematics of the process.  Indeed, the HB cross section (contrary to the SB one) is by itself, at least in principle, a physically measurable quantity for any kinematically allowed value of this threshold. All calculations have been carried out without neglecting the electron mass.   The paper is organized as follows. In sect.~2 we discuss the lowest order prediction for the neutrino electron differential cross section, together with its soft bremsstrahlung and one--loop QED corrections. The hard bremsstrahlung recoil electron spectrum is examined in sect.~3. In the same section, this contribution is added to the virtual and soft ones to derive the total QED corrections.  In sect.~4 we evaluate the spectrum of the total combined energy of the recoil electron and a possible accompanying photon emitted in the scattering process. We summarize our main results in sect.~5, discussing their applicability and emphasizing their role in the analysis of solar neutrino electron scattering experiments. ",
        "conclusions": "When are the results of sects.~2, 3 and 4 applicable? In sects.~2 and 3 we evaluated the $O(\\alpha)$ SM prediction for the electron spectrum in the reaction $\\nu_l + e \\rightarrow \\nu_l + e \\;(+\\gamma)$ (eq.~(\\ref{eq:SMdE})), where $(+\\gamma)$ indicates the possible emission of a photon. In this calculation we assumed that the final--state photon is not detected and, as a consequence, we integrated over all possible values of the photon energy $\\omega$. Therefore, eq.~(\\ref{eq:SMdE}) is the appropriate theoretical prediction to use in the analysis of $\\nu$--$e$ scattering when the detector is completely blind to photons of all energies, but can precisely measure $E$, the energy of the electron. Of course, a detector could provide more information by detecting photons as soon as their energy is above an experimental threshold $\\epsilon$.  In this case, still assuming a precise determination of $E$, one can employ eq.~(\\ref{eq:SMdE}), minus its HB correction, to analyze those events which are counted as nonradiative (elastic), while the HB part can be used, at least in principle, for a separate determination of the inelastic cross section. Indeed, contrary to previous calculations, our predictions are valid for an arbitrary value of the threshold $\\epsilon$ (and include the previously unknown $LR$ term).    In sect.~4 we evaluated the spectrum of the total combined energy of the recoil electron and a possible accompanying photon emitted in the scattering process (eq.~(\\ref{eq:SMdeom}); $\\eom$ was defined in eq.~(\\ref{eq:eom})).  This type of analysis is useful when the photon energy $\\omega$ cannot be separately determined although it fully contributes to the precise total energy measurement if its value is above a specific threshold $\\epsilon$. Let's consider an experimental setup able to measure the photon energy if it's higher than $\\epsilon$, but completely blind to low energy photons $(\\omega\\! <\\!\\epsilon)$.  Let's also assume that the electron energy $E$ is precisely measurable independently of its value. This detector can determine both the differential cross section $d\\sigma/dE_\\omega$ (eq.~(\\ref{eq:SMdeom})) and the electron spectrum $d\\sigma/dE$ (eq.~(\\ref{eq:SMdE})) (as well as its separate HB component).  There are experiments, however, which cannot measure $E$, but only $\\eom$, with a specific value of the threshold $\\epsilon$. BOREXINO \\cite{BX} and KamLAND \\cite{KL}, for example, are liquid scintillation detectors in which photons and electrons induce practically the same response. If a photon is emitted in the $\\nu$--$e$ scattering process, its energy $\\omega$ is counted together with $E$, provided their sum lies within a specific range. The appropriate theoretical prediction for their analysis is given, therefore, by the cross section $d\\sigma/dE_\\omega$ of eq.~(\\ref{eq:SMdeom}) with a very small value of $\\epsilon$ (for the case $\\nu=0.862\\mev$ see the thin lines in fig.~8). However, we should point out that although the QED corrections $\\fx(E)$ (in eq.~(\\ref{eq:SMdE})) and $\\fxov(\\eom,\\epsilon)$ with small $\\epsilon$ (in eq.~(\\ref{eq:SMdeom})) are different, their numerical values are very small when $\\nu=0.862\\mev$, the energy of the monochromatic neutrinos produced by electron capture on \\Be7 in the solar interior.  In fact, as shown in fig.~8, both $(\\alpha/\\pi)\\fx(E)$ and $(\\alpha/\\pi)\\fxov(\\eom,\\epsilon)$ with small $\\epsilon$ are in this case of $O(\\lsim 1\\%)$, and neither of the above collaborations is likely to reach this high level of accuracy in their analyses of the crucial \\Be7 line.   There are detectors in which it might not be possible to identify the measured energy with either $E$ or $\\eom$. Indeed, the electron and the photon may produce indistinguishable signals and the total observed energy might not be the simple sum of $E$ and $\\omega$, but some other function of these two variables. Super--Kamiokande (SK)\\cite{SK}, for example, a water Cherenkov counter measuring the light emitted by electrons recoiling from neutrino scattering, uses the number of hit photomultiplier tubes to determine the electron energy. However, a photon emitted in the scattering process may induce additional hits indistinguishable from those of the electron. Moreover, a photon and an electron of the same energy may produce different numbers of hits and, therefore, it might not be possible to identify the total measured energy with the sum $E+\\omega$.   SK measures solar neutrinos with energies varying from 5 to 18 MeV. For $\\nu=10\\mev$, fig.~9 shows that the QED corrections to the differential cross sections $d\\sigma/dE$ (eq.~(\\ref{eq:SMdE})) and $d\\sigma/dE_\\omega$ (eq.~(\\ref{eq:SMdeom})) are of $O(1\\%)$. Corrections of this order may be relevant for the analysis of the very precise data obtained by this collaboration. In fact, SK's Monte Carlo simulations of the expected energy spectrum of recoil electrons from solar neutrino scattering include the QED corrections of ref.~\\cite{BKS} (as well as the EW ones). As we investigated in sect.~3, these corrections provide good approximations of the complete $O(\\alpha)$ QED corrections $\\fx(E)$ to the electron spectrum of eq.~(\\ref{eq:SMdE}) (see fig.~4). Our previous discussion, however, seems to suggest that these corrections may not be appropriate for SK's solar neutrino analysis.  On the other hand, the SM prediction for the spectrum of the combined energy of electron and photon of sect.~4 (eq.~(\\ref{eq:SMdeom})) may be suitable, but only if we can assume a similar efficiency in the detection of photons and electrons, and if also relatively low energy electrons contribute to the total energy measurement.  If these conditions are not met, and the precision of the data requires it, one should probably perform a dedicated analysis of the double differential cross section $d^2\\sigma/(dE \\,d\\omega)$ with a response function specifically designed for this detector. A triple differential cross section $d^3\\sigma/(dE \\,d\\omega \\,d\\phi)$, where $\\phi$ is the angle between the directions of the electron and the photon, may also be useful (see the third article of ref.~\\cite{ZKN}, and ref.~\\cite{BBBS}).   We will conclude by noticing that the QED corrections, contrary to the EW ones, depend strongly on the initial neutrino energy and become sizeable at high energies. Their appropriate expression will have to be taken into account in the analysis of future precise high energy $\\nu$--$e$ scattering experiments.      \\vspace{0.5cm}\\begin{center}{\\large\\bf Acknowledgments}\\end{center}  \\noindent I would like to thank Prof.~J.~N.~Bahcall and Prof.~A.~Sirlin for suggesting this problem and for enlightening discussions. I am particularly indebted to Prof.~P.~Minkowski and Prof.~A.~Sirlin for reading the manuscript and for very useful comments. I would also wish to express my gratitude to J.~Arafune, D.~Bardin, C.~Greub, A.~Held, K.~Holland, T.~Kinoshita, W.~Marciano, M.~Nakahata, J.~Papavassiliou, M.~Samaras, M.~Schaden, S.~Schoenert, Y.~Takeuchi and Y.~Totsuka for instructive conversations and exchanges of ideas. This work was supported by Schweizerischer Nationalfonds."
    },
    "0011/astro-ph0011285_arXiv.txt": {
        "abstract": "We present the results of a quantitative study of the degree of extension to the boundary of the classical convective core within intermediate mass stars. The basis of our empirical study is the stellar population of four young populous clusters in the Magellanic Clouds which has been detailed in Keller, Bessell \\& Da Costa (\\cite{kel00}). The sample affords a meaningful comparison with theoretical scenarios with varying degrees of convective core overshoot and binary star fraction. Two critical properties of the population, the main-sequence luminosity function and the number of evolved stars, form the basis of our comparison between the observed data set and that simulated from the stellar evolutionary models. On the basis of this comparison we conclude that the case of no convective core overshoot is excluded at a 2 $\\sigma$ level. ",
        "introduction": "The young populous clusters of the Magellanic Clouds (MCs) are amongst the richest, resolvable young clusters in the Local Group. The clusters provide a statistically significant sample of stars with which to constrain the evolutionary behaviour of massive stars. This paper is the second in a series detailing our study of four young populous clusters, NGC 330, 1818, 2004 \\& 2100. These clusters have ages of 15-50 Myrs and main-sequence terminus masses in the range of 9-12$\\rm{M}_{\\odot}$. In this mass range that the treatment of convection, particularly within the convective core, is critical to the course of stellar evolution. Ongoing debate focuses on the degree of extension of the convective core beyond that predicted by standard, non-rotating stellar evolutionary models. This paper seeks to provide quantitative constraints on the magnitude of convective core extension through confrontation of our observations with the predictions of stellar evolutionary models. Details of our observations are given in (Keller, Bessell \\& Da Costa 2000, hereafter Paper 1).  Extension of the convective core has traditionally been discussed in terms of the formalism of convective core overshoot. Convective core overshoot (CCO) is one easily visualised (and parameterised!) mechanism for extending the size of the convective core. Under this description, gas packets rising in convective motion from the core reach a distance from the centre where the acceleration imparted to the convective element by the buoyancy force vanishes. This height defines the classical Schwarzschild convective core boundary. Such packets arrive at this height with non-zero velocity, hence some degree of overshoot must occur. However, the classical picture considers that this degree of overshoot is negligible (Saslaw \\& Schwarzschild \\cite{sas65}).  Numerous studies have attempted to ascertain the efficiency of CCO from a theoretical basis with different results, ranging from negligible to substantial (see e.g. Bressan et al.\\ \\cite{bre81}). An analytical approach appears limited given the complexity of the phenomenon. Laboratory fluid dynamics shows that an understanding of mixing requires a description of the turbulence field at all scales, a problem which will require detailed hydrodynamical modelling.  Recourse to observation of stellar populations is required to ascertain the amount of convective core extension to apply within stellar evolutionary models. Demonstration that the convective core is of greater extent than the classical core does not, however, determine its causation. Several effects can bring about an extension of the convective core above that predicted by classical models. Opacity, for instance is critical in determination of the convective core size. The superseding of the old LAOL opacities (Huebner et al.\\ \\cite{hub77}) by those of OPAL (Iglesias et al. \\cite{igl92}) resulted in a substantial increase in the size of the classical convective core (see Stothers \\& Chin \\cite{sto92}). Further, recent studies indicate that rotation can provide a natural way to bring about increased internal mixing and hence a larger core (Heger \\& Langer \\cite{heg00} and Meynet \\& Maeder \\cite{mey00}). Nevertheless, the formalism of CCO offers a straightforward parameterisation of the effect of convective core extension regardless of causation. For this reason the following discussion is in terms of CCO.  CCO brings about a number of important evolutionary changes which are expressed in the colour-magnitude diagram (CMD) of a cluster population. In particular, the MS lifetime is extended and stars evolve on the main-sequence (MS) to cooler temperatures and higher luminosities relative to non-overshoot models. Further, the subsequent post-MS evolution occurs at a more rapid pace, leading to a number ratio of He burning stars to H burning stars that is strongly dependent on the degree of overshoot ($\\Lambda_{c}$; Langer \\& El Eid \\cite{lan86}, Bertelli et al.\\ \\cite{ber86}).  Through comparison of the observed cluster CMD with the predictions of stellar evolutionary models containing various degrees of CCO we are able to ascertain the amount of overshoot exhibited by the population.  In considering massive stars, however, it is necessary to examine young clusters, which in the case of our Galaxy are too poorly populated to individually yield much insight. It is possible to combine data from sets of similar aged clusters as in the work of Mermilliod \\& Maeder (\\cite{mer86}) and Maeder \\& Mermilliod (\\cite{mae81}), yet such sets are inevitably heterogeneous in age and metallicity. The young populous clusters of the Magellanic Clouds contain upward of 10$\\times$ the mass of their Galactic counterparts and thus provide ideal targets for study.  Perhaps the best studied system is NGC 1866 (age=100-200Myr; Testa et al. \\cite{tes99}). The intermediate age population of this system has been considered the proof both for (Becker \\& Mathews \\cite{bec83}, Chiosi et al.\\ \\cite{chi89}) and against (Brocato et al.\\ \\cite{bro94}, Testa et al.\\ \\cite{tes99}) CCO. In Chiosi et al.\\ (ibid) the degree of CCO exhibited by the population on NGC 1866 is estimated at $\\Lambda_{c}$=0.5. Whilst remaining a contested issue, this degree of overshoot then forms the basis of the physical input supplied to the standard evolutionary models (SEMs) of the Padua group (e.g. Fagotto et al.\\ \\cite{fag94}). The motivation for the present study is to extend the number of young clusters which have been examined in regard to the amount of CCO and to place the need or otherwise for CCO on a firmer basis.  Previous ground-based and IUE studies of the clusters of the present work have revealed several features which are anomalous compared to SEMs which do not include overshooting. These features are indicative of some degree of convective overshoot within the population. However, further conclusions from these studies are limited since the observations have been restricted to the brightest members, the only stars for which accurate effective temperatures were attainable. They extend to just below the MS turnoff. In addition previous studies have made use of reddening, metallicities, and in some cases opacities, which have since been superseded by more appropriate input. The present paper revisits and forms a significant extension to these previous studies. We draw together new data presented in Paper 1 together with a set of stellar evolutionary models containing homogeneous input physics. ",
        "conclusions": "In the analysis to follow we will focus of the MS integrated luminosity function (MSILF) for the purpose of comparing model with data. The implications of the effective temperature information is examined separately in Section \\ref{efftemp}. The MSILF is formed for each magnitude ($V_{bin}$) by the summation of the number of objects with $V<V_{bin}$ (top panel of Figure~\\ref{figsig0.00}). The comparison between Model \\#1, incorporating no extension of the convective core, and NGC 330 is summarised in Figure~\\ref{figsig0.00} for three ages: at log age=7.30, 7.36 and 7.50. In order to examine the difference between the MSILF observed and that simulated, we have constructed the quantity $Q=(\\rm{N_{obs}-N_{theo}}) / (\\sqrt{\\rm{N_{theo}} + \\rm{N_{obs}}})$ (seen in the bottom panel of Figure~\\ref{figsig0.00}).  $Q$ is constructed in this manner in order to provide a clear indication of where significant deviations occur in the model predictions. The nature of the problem means that these deviations occur towards the top of the MS, for this reason other well known statistical quantifiers such as the K-S statistic are not well suited to its description.  As the total number of stars brighter than $\\rm{M}_{V}<$$V_{min}$ is equal to the observed number, $V_{min}$ by definition corresponds to $Q$=0. Moving towards brighter magnitudes, the value of $Q$ will not necessarily remain around zero: the larger the range over which $Q$ remains close to zero, the better the fit.  Consider the case of log age=7.30; here the MS terminus point is in close agreement with that observed: however, the model does not mimic the observed {\\it{density}} on the MSILF. The model produces too high a density towards the end of the MS, hence the positive values of $Q$. In the case of a log age=7.5, the model fails to reproduce the MS terminus, which results in a region of negative $Q$. The area, $A=\\sum |Q|$, in the above figures gives a global description of the goodness of fit for each model to the data. In the case of Model \\#1 the maximum likelihood is obtained for a log age of 7.36.  \\placefigure{figsig0.00}  We have performed a quadratic interpolation between the three models in order to sample the range of CCO from 0.0 to 1.0. This is justified as the effects brought about by modification of the CCO level are well behaved. We have repeated the determination of $A$ 100 times for each value of age to enable us to describe the stochastic uncertainty involved. We repeat for values of CCO in the range covered by our models. In the parameter space of each cluster there exists a minimum, $A_{min}$, which represents the favoured model. We express every other determination of $A$ in terms of the number of standard deviations from $A_{min}$. The normalised probability density at each point is then given by $L$=$A_{min}e^{-\\sigma^{2}/2}$. The resulting probability density is shown in Figures~\\ref{lfs} as a greyscale plot with a 1$\\sigma$ contour level overlaid.  We find that this method provides a more robust determination of the age of the cluster than is available by a simple isochrone fit to the data. This method makes use of the bulk of essentially unevolved MS members whose ILF depends significantly on the assumed age. On the other hand the degree of extension to the convective core is only weakly constrained, but in the case of each cluster favours a non-zero level of CCO.  It is important to consider why the MSILF of these clusters is best fit by a significant level of CCO. A comparison of the best-fitting models to the MSILF of NGC 330 is shown in Figure~\\ref{compbestfits}. The differences between the models of various levels of overshoot are due to a change in the density along the MS. Models with a larger degree of overshoot evolve towards higher luminosities in the final stages of MS evolution which results in a stretching out of the MSILF towards the MS terminus. The lower panel of Figure~\\ref{compbestfits} shows this clearly. The $\\Lambda_{c}$=0.0 model terminates too rapidly to mimic the observations, the $\\Lambda_{c}$=1.0 model on the other hand terminates too slowly.  \\placefigure{lfs}  \\placefigure{compbestfits}    Our study of four young populous clusters in the Magellanic Clouds has sought to quantify the extension to the classical Schwarzschild stellar convective core by convective core overshoot (CCO). We have compared the H-R diagrams of each cluster with synthetic ones produced from a set of stellar evolutionary models different only in their degree of CCO.  Using the twin constraints of (1) the main-sequence LF and (2) the number of evolved stars, we locate the age and CCO range which maximises the agreement between theory and observation. In each cluster we find that the case of no-overshoot is excluded to a 2$\\sigma$ level. With the assumption of a Salpeter mass function and a 30\\% binary fraction within the model population, an average of the results for each cluster yields an overshoot parameter $\\Lambda_{c}$=0.62$\\pm$0.22. A negligable binary fraction favours $\\Lambda_{c}$=0.56$\\pm$0.19. Both results are formally in agreement with the amount of CCO included within the widely used stellar evolutionary models of the Padua and Geneva groups."
    },
    "0011/astro-ph0011550_arXiv.txt": {
        "abstract": "We do not fully understand the dynamics and evolution of a radio galaxy.  Models of classical double (Type II) sources are in a reasonable state, but these objects are rare.  Non-type II sources (generically called Type I) are far more common, but much less well understood. In this paper I use the data to suggest possible new models for Type I sources, and discuss the physical questions which these new models raise. ",
        "introduction": "What is a Type I radio source?  This question can be answered in terms of the data (what does it look like?), or in terms of the dynamics (how does it work?).  In neither case is the answer clear.  The observational answer has changed over the years.  The initial Fanaroff-Riley definition combined morphology (inner hot spots, close to the galaxy) with low radio power ($P_{rad}$). Later, Ledlow and Owen (1996) complicated the issue by showing that the division in $P_{rad}$ between Types I and II depends on the luminosity of the parent galaxy, $L_{opt}$.  Now, with the advent of large radio samples, it has become clear that the morphology issue is also more complicated than was initially realized.  When we study images of radio galaxies which do not have strong outer hot spots (that is, which would not be called Type II), we soon discover these sources are not a uniform set.  The dynamical definition of a Type I is equally unclear. We can pose several questions which should be answered by a dynamical model of radio sources.  How does a given source evolve?  What governs its size, morphology, and radio power?  How do these quantities change with age?  We have a reasonable picture for Type II sources, which in principle allows quantitative answers to these questions.  We do not, however, have a similarly useful picture for non-Type II sources, even though they constitute most of the radio source population.  This is my concern in this paper. As a shorthand, I will refer to any source which is not clearly a Type II, as a Type I.   In what follows I will discuss  the morphology and dynamics of these interesting sources, combining the material from my talk and poster. ",
        "conclusions": "All  I can really conclude is that Type I sources remain a fascinating challenge to our understanding of radio galaxy evolution. I anticipate that combining high-quality imaging with new theoretical analysis will improve our understanding,  and no doubt expose the naivety of my speculations here."
    },
    "0011/astro-ph0011416_arXiv.txt": {
        "abstract": "The role of Horizontal Showers induced by cosmic rays is discussed in detail. A new approach to the calculation of the muon component in horizontal air showers induced by protons, heavier nuclei or photons is presented which allows a simple analytical evaluation of the muon density profiles at ground level. The results of the first application of these results to horizontal air showers detected at the Haverah Park Array by the recently started {\\sl Leeds-Santiago collaboration}, leading to important restrictions on composition at ultra high energies, are reported. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011167_arXiv.txt": {
        "abstract": "{{\\it XMM-Newton} Reflection Grating Spectrometer (RGS) spectra of the Narrow Line Seyfert 1 galaxies \\object{MCG~$-$6-30-15} and \\object{Mrk~766} are physically and spectroscopically inconsistent with standard models comprising a power-law continuum absorbed by either cold or ionized matter.  We propose that the remarkably similar features detected in both objects in the 5 -- 35 \\AA\\ band are H-like oxygen, nitrogen, and carbon emission lines, gravitationally redshifted and broadened by relativistic effects in the vicinity of a Kerr black hole.  We discuss the implications of our interpretation, and demonstrate that the derived parameters can be physically self-consistent. ",
        "introduction": "The precise shape of the low energy spectra of active galaxies has traditionally been very difficult to establish.  The combined effects of interstellar absorption, moderate spectral resolution of available detectors, and intrinsic complexity in the sources have so far prevented us from determining even whether the underlying spectrum is mainly due to continuum emission, or includes discrete emission and absorption components. The model generally adopted to match the observations is that of a continuum spectrum absorbed by partially ionized material (Halpern \\cite{halpern84}; Reynolds \\cite{reynolds97}, and references therein); the origin and location of this warm absorber at the core of active galaxies, however, is still very much a matter of debate (e.g., Otani et al.\\ \\cite{otani96}).  The enhancement in energy resolution and sensitivity afforded by the {\\it XMM-Newton} Reflection Grating Spectrometer (RGS; den Herder et al.\\ \\cite{denherder01}) provides us with the potential to unravel the true origin of the soft X-ray emission in AGN for the first time. RGS observations of \\object{MCG~$-$6-30-15} and \\object{Mrk~766}, which are reported here, have forced us to examine alternatives to the warm absorber model, and to propose a new and radically different interpretation of the soft X-ray spectra of active galaxies. ",
        "conclusions": "We have shown that a simple warm-absorber interpretation of the RGS spectra of \\object{MCG~$-$6-30-15} and \\object{Mrk~766} is untenable on spectroscopic grounds.  Broad line emission from a relativistic disk surrounding a maximally rotating Kerr black hole seems to explain the data remarkably well.  The physical self-consistency of this scenario remains to be established, however, the preliminary analysis presented in Sec. 4 is encouraging.  Note that the conclusions we draw do not depend on any pre-conceived assumption about the shape of the ionizing continuum.  This result could not have been achieved without the combination of large effective area and high energy resolution afforded by the {\\it XMM-Newton} RGS.  The poorer resolution of CCD spectrometers cannot provide discriminatory power for the warm absorber versus line emission debate raised by the RGS results presented in this paper.  A more robust test will be finding the same problems and applying the same solution to other AGN sources."
    },
    "0011/astro-ph0011484_arXiv.txt": {
        "abstract": "Corresponding to the recent observational claims that we are in a local void (an underdense region) on scales of $200  - 300$ Mpc, the magnitude-redshift relation in a cosmological model with a local void is investigated. It is already evident that the accelerating behavior of high-z supernovae can be explained in this model, because the local void plays a role similar to the positive cosmological constant. In this paper the dependence of the behavior on the gaps of cosmological parameters in the inner (low-density) region and the outer (high-density) region, the radius of the local void, and the clumpiness parameter is studied and its implication is discussed. ",
        "introduction": "One of the most important cosmological observations at present is the [$m, z$] relation for high-z supernovae (SNIa), which play a role of standard candles at the stage reaching epochs $z \\ga 1$. So far the observed data of SNIa have been compared with the theoretical relation in homogeneous and isotropic models, and many workers have made efforts to determine their model parameters (Garnavich et al. 1998, Schmidt et al. 1998, Perlmutter et al. 1999, Riess et al. 1998, 2000, Riess 2000).  Here is, however, an essentially important problem to be taken into consideration. It is the homogeneity of the Universe. According to Giovanelli et al. (1998, 1999) and Dale et al. (1999) the Universe is homogeneous in the region within $70 h^{-1}$ Mpc \\ (the Hubble constant $H_0$ is $100 h$ km s$^{-1}$ Mpc$^{-1}$). On the other hand, recent galactic redshift surveys (Marinoni et al. 1999, Marzke et al. 1998, Folkes et al. 1999, Zucca et al. 1997) show that in the region around $200 - 300 h^{-1}$ Mpc from us the distribution of galaxies may be inhomogeneous. This is because the galactic number density in the region of $z < 0.1$ or $< 300 h^{-1}$ Mpc from us was shown to be by a factor $> 1.5$ smaller than that in the remote region of $z > 0.1$. Recently a large-scale inhomogeneity suggesting a wall around the void on scales of $\\sim 250 h^{-1}$ Mpc has been found by Blanton et al. (2000) in the SDSS commissioning data (cf. their Figs. 7 and 8). Similar walls on scales of $\\sim 250 h^{-1}$ Mpc have already been found in the Las Campanas and 2dF redshift surveys near the Northern and Southern Galactic Caps (Shectman et al. 1996, Folkes et al. 1999, Cole et al. 2000). These results mean that there is a local void with the radius of $200 - 300 h^{-1}$ Mpc and we live in it.  Moreover, the measurements by Hudson et al. (1999) and Willick (1999) for a systematic deviation of clusters' motions from the global Hubble flow may show some inhomogeneity on scales larger than $100 h^{-1}$ Mpc. Another suggestion for inhomogeneity comes from the periodic wall structures on scales of $\\sim 130 h^{-1}$ Mpc, as have been shown by Broadhurst et al. (1990), Landy et al. (1996), and Einasto et al. (1997). This is connected with the anomaly of the power spectrum around $100 - 200 h^{-1}$ Mpc (so-called ``excess power'') which was discussed by Einasto et al. (1999). This fact also may suggest some inhomogeneity in the above nearby region.  If the local void exists really, the Hubble constants also must be inhomogeneous, as well as the density parameters, and the theoretical relations between observed quantities are different from those in homogeneous models. At present, however, the large-scale inhomogeneity of the Hubble constant has not been observationally established yet because of the large error bars in the various measurements (cf. Tomita 2001).  In my previous papers (Tomita 2000a and 2000b) cited as Paper 1 and Paper 2, I showed various models with a local void and discussed the bulk flow, CMB dipole anisotropy, distances and the [$m, z$] relation in them in the limited parameter range. It was found that the accelerating behavior of supernovae can be explained in these models without cosmological constant. On the other hand, Kim et al. (1997) showed that the difference between the local and global values of the Hubble constant should be smaller than $10 \\%$ in {\\it homogeneous} cosmological models in order to be consistent with the SNIa data. However, this does not impose any strong condition on the difference in inhomogeneous models, because their analyses were done using the luminosity distance in {\\it homogeneous} models and so they are incomplete.  In fact my previous papers showed concretely that in inhomogeneous models larger differences can be consistent with the data. The possibility that the above difference may explain the behavior of SNIa was later discussed also by Goodwin et al. (1999).  In this paper I describe first (in \\S 2) a simplified cosmological model with a local void and treat distances in light paths with nonzero clumpiness (smoothness) parameter $\\alpha$. In the previous paper (Paper 2), I considered only distances in full-beam light paths ($\\alpha = 1$), but in the realistic paths there are deviations from $\\alpha = 1$ due to lensing effects from inhomogeneous matter distributions. In \\S 3, I show the dependence of the [$m, z$] relation on model parameters such as the radius of the local void, the ratios of density parameters and Hubble constants in the inner (low-density) and outer (high-density) regions, and the clumpiness parameter. The constraints to the parameters are derived in comparison between the above relations in the present models and the relations in homogeneous models.  Finally (in \\S 4), we discuss the remaining problems and describe concluding remarks.  \\section[]{Distances in models with a local void}  The inhomogeneous models we consider consist of inner (low-density) region V$^{\\rm I}$ and outer (high-density) region V$^{\\rm II}$ which are separated by a single shell. It is treated as a spherical singular shell and the mass in it compensates the mass deficiency in V$^{\\rm I}$. So V$^{\\rm I}$ and the shell are regarded as a local void and the wall, respectively. The line-element in the two regions are \\begin{eqnarray} \\label{eq:m1} ds^2 &=& g^j_{\\mu\\nu} (dx^j)^\\mu (dx^j)^\\nu \\cr &=& - c^2 (dt^j)^2 + [a^j (t^j)]^2 \\Big\\{ d (\\chi^j)^2 + [f^j (\\chi^j)]^2 d\\Omega^2 \\Big\\}, \\end{eqnarray} where $j \\ (=$ I or II) represents the regions, $f^j (\\chi^j) = \\sin \\chi^j, \\chi^j$ and $\\sinh \\chi^j$ for $k^j = 1, 0, -1$, respectively, and $d\\Omega^2 = d\\theta^2 + \\sin^2 \\theta  d\\varphi^2$. In the following the negative curvature is assumed in all regions. The Hubble constants and density parameters are expressed as $(H_0^{\\rm I}, H_0^{\\rm II})$ and $(\\Omega_0^{\\rm I}, \\Omega_0^{\\rm II})$, where we assume that $H_0^{\\rm I} > H_0^{\\rm II}$ and $\\Omega_0^{\\rm I} < \\Omega_0^{\\rm II}$. The distances of the shell and the observer O (in V$^{\\rm I}$) from the centre C (in V$^{\\rm I}$) are assumed to be $200$ and $40 h_{\\rm I}^{-1}$ as a standard case. This shell corresponds to the redshift $\\bar{z}_1 = 0.067$ \\ (see Figure 1).  \\begin{figure} \\centerline{\\psfig{figure=f01.eps,width=\\plotwidth}} \\caption{Model with a spherical single shell. Redshifts for observers at O and C are $z$ and $\\bar{z}$. \\label{fig:1}} \\end{figure}  In Paper 2 we derived the full-beam distances ($\\overline{\\rm CS}$) between the centre C and a source S, and the distances ($\\overline{\\rm OS}$) between an observer O and S. The two distances are nearly equal in the case when $\\overline{\\rm CS}$ or $\\overline{\\rm OS}$ is much larger than $\\overline{\\rm CO}$. Since we notice this case alone in the following, we treat the light paths $\\overline{\\rm CS}$ for simplicity. Then the angular-diameter distance $d_{\\rm A}$ is \\begin{equation} \\label{eq:m2} d_{\\rm A} = a^{\\rm I}(\\bar{\\eta}^{\\rm I}_{\\rm s}) \\sinh (\\bar{\\chi}^{\\rm I}_{\\rm s}), \\end{equation} if a source S is in V$^{\\rm I}$, where $(\\bar{\\eta}^{\\rm I}_{\\rm s}, \\bar{\\chi}^{\\rm I}_{\\rm s})$ are the coordinates of S, and $\\eta$ is the conformal time coordinate. Here bars are used for the coordinates along the light paths to the virtual observer at C. If S is in V$^{\\rm II}$, we have \\begin{eqnarray} \\label{eq:m3} d_{\\rm A} = a^{\\rm I}(\\bar{\\eta}^{\\rm I}_1) \\sinh (\\bar{\\chi}^{\\rm I}_1) &+& [a^{\\rm II}(\\bar{\\eta}^{\\rm II}_{\\rm s} \\sinh (\\bar{\\chi}^{\\rm II}_{\\rm s})\\cr &-& a^{\\rm II}(\\bar{\\eta}^{\\rm II}_1) \\sinh (\\bar{\\chi}^{\\rm II}_1)], \\end{eqnarray} where $(\\bar{\\eta}^{\\rm I}_{\\rm 1}, \\bar{\\chi}^{\\rm I}_{\\rm 1})$ stand for the shell, and we have \\begin{equation} \\label{eq:m4} a^{\\rm I}(\\bar{\\eta}^{\\rm I}_1) \\sinh (\\bar{\\chi}^{\\rm I}_1) = a^{\\rm II}(\\bar{\\eta}^{\\rm II}_1) \\sinh (\\bar{\\chi}^{\\rm II}_1) \\end{equation} from the junction condition.  Here we treat the following equation for the angular-diameter distance to consider the clumpiness along paths into account (Dyer and Roeder 1973, Schneider et al. 1992, Kantowski 1998, Tomita 1999): \\begin{eqnarray} \\label{eq:m5} {d^2 (d_{\\rm A}^j) \\over d(z^j)^2} &+& \\Big\\{{2 \\over 1+z^j}+{1 \\over 2}(1+z^j) \\Big[\\Omega_0^j (1+3z^j) + 2 -2\\lambda_0^j\\Big]\\cr && \\times F^{-1} \\Big\\}{d (d_{\\rm A}^j) \\over dz^j} + {3 \\over 2} \\Omega_0^j \\alpha (1+z^j) F^{-1} d_{\\rm A}^j = 0, \\end{eqnarray} where $j = $ I and II, $z^j$ is the redshift in the region V$^j$, $\\alpha$ is the clumpiness parameter, and \\begin{equation} \\label{eq:m6} F \\equiv (1+ \\Omega_0^j z^j)(1+z^j)^2 - \\lambda_0^j z^j (2+ z^j). \\end{equation} Here and in the following the bars are omitted for simplicity. The two redshifts at the shell are equal, i.e. \\begin{equation} \\label{eq:m7} z_1^{\\rm I} = z_1^{\\rm II} \\ (\\equiv z_1) \\end{equation} for the comoving shell (cf. Paper I).  The distances $d_{\\rm A}^{\\rm I}$ in V$^{\\rm I}$ is obtained solving Eq.(\\ref{eq:m5}) under the conditions at $z^{\\rm I} = 0$: \\begin{equation} \\label{eq:m8} (d_{\\rm A}^{\\rm I})_0 = 0, \\quad (d_{\\rm A}^{\\rm I}/dz^{\\rm I})_0 = c/H_0^{\\rm I}, \\end{equation} and $d_{\\rm A}^{\\rm II}$ in V$^{\\rm II}$ is obtained similarly under the conditions at $z^{\\rm II} = 0$: \\begin{equation} \\label{eq:m9} (d_{\\rm A}^{\\rm II})_0 = const, \\quad (d_{\\rm A}^{\\rm II}/dz^{\\rm II})_0 = c/H_0^{\\rm II}, \\end{equation} where $const$ is determined so that the junction condition $d_{\\rm A}^{\\rm I}(z_1) = d_{\\rm A}^{\\rm II}(z_1)$ may be satisfied at the shell. Then the distance $d_{\\rm A} (z_s)$ from C to the source S is \\begin{equation} \\label{eq:m10} d_{\\rm A} (z_s) = d_{\\rm A}^{\\rm I}(z_s) \\hspace{3cm} {\\rm for} \\ z_s \\leq z_1, \\end{equation} and \\begin{equation} \\label{eq:m11} d_{\\rm A} (z_s) = d_{\\rm A}^{\\rm I}(z_1) + d_{\\rm A}^{\\rm II}(z_s) - d_{\\rm A}^{\\rm II}(z_1) \\quad {\\rm for} \\ z_s > z_1, \\end{equation} where $z_s = z_s^{\\rm I}$ and $z_s^{\\rm II}$ for $z_s \\leq z_1$ and $z_s > z_1$, respectively. The luminosity distance $d_{\\rm L}$ is related to the angular-diameter distance $d_{\\rm A}$ \\ by \\ $d_{\\rm L} = (1 + z)^2 d_{\\rm A}$.  As for the clumpiness parameter $\\alpha$, we studied the distribution function $N(\\alpha)$ as a function of $z$ in our previous papers (Tomita 1998, Tomita 1999).  To obtain  $N(\\alpha)$, we first derived model universes consisting of galaxies and halos using a N-body simulation technique, secondly calculated the angular-diameter distance by solving null-geodesic equations along many light paths between an observer and sources at epoch $z$, and finally derived a statistical distribution of $\\alpha$ determined in a comparison with the Friedmann distance ($\\alpha = 1.0$) and the Dyer-Roeder distance ($\\alpha = 0.0$). As the result of these studies, it was found that the average value $\\bar{\\alpha}$ of $\\alpha$ is $1.0$, which represents the Friedmann distance, and the dispersion $\\sigma_\\alpha$ can be $\\sim 0.5$ for $z <2.0$. If the detection of high-z supernovae is done in completely random directions, the observed average value of  $\\alpha$ is equal to the above theoretical average value $\\bar{\\alpha}$. But if the detections are biased to the directions with less galactic number per steradian to avoid the dust obscuration, we may have the value of $\\alpha \\sim \\bar{\\alpha} - \\sigma_\\alpha$. Then the angular-diameter and luminosity distances are somewhat longer than the average Friedmann distances. In the next section we show the cases with $\\alpha = 1.0, 0.5,$ and $0.0$ for comparison. The lensing effect on the [$m, z$] relation of SNIa was discussed also by Holz (1998), Porciani and Madau (2000), and Barber (2000).  \\begin{figure} \\centerline{\\psfig{figure=f02.ps,width=\\plotwidth}} \\caption{The [$m, z$] relation in cosmological models with a local void. The solid line denotes the case with a standard parameter set given in (\\ref{eq:p1}). The dotted and dash lines stand for homogeneous models with $(\\Omega_0, \\lambda_0) = (0.3, 0.7)$ and $(0.3, 0.0)$, respectively, for comparison. \\label{fig:2}} \\end{figure} ",
        "conclusions": "As for the [$m, z$] relation in cosmological models with a local void, we studied the parameter dependence of their accelerating behavior, and found that the local void with $r_1 \\la 200 h^{-1}$ Mpc, $H_0^{\\rm II}/H_0^{\\rm I} \\ga 0.82$, and $\\Omega_0^{\\rm II} \\la 0.6$ is appropriate for explaining the accelerating behavior of SNIa without cosmological constant, that the lensing with $\\alpha \\sim 0.5$ is effective at epochs of $z \\ga 1.0$, and that the cosmological constant ($\\lambda_0^{\\rm II} \\sim 0.4$) necessary for flatness in the outer region has a role supplementary to the accelerating behavior. On the basis of these results we can determine in the next step what values of the model parameters are best in the direct comparison with the observational data of SNIa.  In the Universe with CDM matter, the probability that the inhomogeneity of Hubble constant $\\delta H/H \\sim 0.2$ on scales $\\sim 200$ Mpc associated with general density perturbations is realized is extremely small, as was clarified and discussed by Turner et al. (1992), Nakamura and Suto (1995), Shi and Turner (1998), and Wang et al. (1998).  The constraint from CMB dipole anisotropy was also discussed by Wang et al.(1998). It should be noticed here that the {\\it spherical} void which we are considering is exceptionally compatible with the constraint from CMB dipole anisotropy, in spite of the above large deviation of Hubble constant, as long as the observers are near the centre (cf. Paper 1). Inversely it may be suggested that the local void on scales $\\sim 200$ Mpc must be {\\it spherical} or {\\it nearly spherical}, if its existence is real.  In comparison with observations of the galactic number count - magnitude relation, on the other hand, Phillips and Turner (1998) have once studied the possibility of an underdense region on scales of $\\sim 300 h^{-1}$ Mpc. However, a necessary wall for the mass compensation has not been considered in their simple models and, in the small-angle observations of the above relation the boundary between the inside region and the outside region was rather vague in contrast to the large-angle redshift surveys.  In the near future the void structure on scales of $\\sim 200 h^{-1}$ Mpc will be clarified by the galactic redshift survey of SDSS in the dominant part of whole sky. Then, observational cosmology will be developed taking into account that we are in a local void.  The author thanks K. Shimasaku for discussions on galactic redshift surveys and V. M$\\ddot{\\rm u}$ller for helpful comments. This work was supported by Grant-in Aid for Scientific Research (No.~12440063) from the Ministry of Education, Science, Sports and Culture, Japan."
    },
    "0011/astro-ph0011021_arXiv.txt": {
        "abstract": "The phases of the Fourier modes appearing in a plane-wave expansion of cosmological density fields play a vital role in determining the morphology of gravitationally-developed clustering. We demonstrate this qualitatively and  quantitatively using simulations. In particular, we use cross-correlation and rank-correlation techniques  to quantify the agreement between a simulated distribution and phase-only reconstructions. The phase-only reconstructions exhibit a high degree of correlation with the original distributions, showing how meaningful spatial reconstruction of cosmological density fields depends  more on phase accuracy than on amplitudes. ",
        "introduction": "The standard theory of the origin of the large-scale structure of the universe involves the assumption that the structure we see today grew by gravitational instability from initial small fluctuations in the density field. In most popular variants of this model, particularly those involving cosmic inflation, the initial fluctuations are of a particularly simple form known as a Gaussian random field. Gaussian random fields are useful because many properties of Gaussian random density fields can be calculated analytically (e.g. Bardeen et al. 1986). Some direct motivation for such an assumption emerges from inflationary models, wherein the fluctuations are generated by quantum oscillations of the scalar field driving inflation (the ``inflation''). Even if inflation turns out to be incorrect, however, the Central Limit Theorem tends to produce Gaussian fluctuations in any linear process, so that they are the most generic form for small initial conditions and a natural default assumption.  One particularly interesting property of Gaussian random fields is that the requirement for the density contrast $\\delta({\\bf x})=[\\rho({\\bf x})-\\rho_{0}]/\\rho_{0}$ to be a Gaussian random field is equivalent to that the real and imaginary parts of its Fourier components $\\tilde{\\delta}_k$, where \\begin{equation} \\delta ({\\bf x}) =\\sum \\tilde{\\delta}({\\bf k}) \\exp(i{\\bf k}\\cdot {\\bf x}), \\end{equation} are independently distributed. In other words, the Fourier modes $\\tilde{\\delta}({\\bf k})$, \\begin{equation} \\tilde{\\delta}({\\bf k})=|\\tilde{\\delta} ({\\bf k})|\\exp(i\\phi_{\\bf_k}), \\label{eq:fourierex} \\end{equation} possess phases $\\phi_{\\bf k}$ which are independently distributed and uniformly random on the interval $[0,2\\pi]$. As the density field is simply a sum over a large number of Fourier modes and if the phases of each of Fourier modes are random, the Central Limit Theorem guarantees  that the resulting superposition of the one-point probability distribution ${\\cal P}(\\delta)$ is close to Gaussian and that all of the field's joint probability distributions are multivariate Gaussians. The statistical properties of an isotropic Gaussian random field are then completely specified by its second-order statistical quantity: the covariance function, or alternatively, its power spectrum $P(k)=\\langle\\tilde{\\delta}^2(k)\\rangle$.  In the framework of gravitational instability, the growth of fluctuations can be understood analytically when the density fluctuation amplitude is small compared to the mean density; the linear perturbation theory tells us that each Fourier mode grows with the same rate independent of wavenumber and the statistical distribution of $\\delta$ remains constant except its variance.  The linear theory breaks down when $\\langle\\delta^2\\rangle$ is comparable with unity or beyond, and different Fourier modes start coupling.  One way to look at the mode coupling is that $\\delta$ is always constrained to value $\\delta \\geq -1$. When the perturbation is small, the tail of  ${\\cal P}(\\delta)$ in the part of $\\delta <-1$ assigned by Gaussian distribution is negligible, because probability of $\\delta<-1$ is small. When  the density field evolves beyond the linear regime, i.e., $\\sigma^2 \\equiv \\langle \\delta^2 \\rangle \\sim 1$, a long tail at high $\\delta$ is generated while lower bound is confined at $\\delta=-1$. Gaussian condition is therefore invalid, in that mode coupling effect causes the initial condition to skew. Terms in the evolution           \\begin{figure} \\caption{Visual demonstration of the importance of phase information for clustering morphology. Plate (a), (c) and (e) have the same phase configuration, so do plate (b) and (d). Plate (a), (d) and (f), however, share the same power spectrum, or alternatively, two-point correlation function, so do (b) and (c). See the text for details. }\\label{demo} \\end{figure} \\newpage   \\noindent equations for the Fourier modes that represent coupling between different modes are of second (or higher) order in $\\delta$ and these are neglected when first-order perturbation theory is considered. Phases of Fourier modes therefore are therefore coupled together in a way which is yet to be fully elucidated but which has been recently investigated by Chiang \\& Coles (2000) and Coles \\& Chiang (2000).   One of the reasons for studying Fourier phases in depth is the question of possible primordial non-Gaussianity of the initial density distribution of the Universe. For example, there have been claims of non-Gaussianity from analysis results in the COBE DMR sky maps of cosmic microwave background using a number of different  diagnostics: bispectrum analysis (Ferreira et al. 1998); the wavelet transform (Hobson et al. 1999); and Minkowski functionals (Schmalzing \\& Gorski 1998). The most direct practical approach, however, is through the distribution of Fourier phases. This diagnostic can also avoid the subtlety of more standard methods that depend on the distribution function ${\\cal P}(\\delta)$(Scherrer et al. 1991), in that a density field with Gaussian single-point density distribution ${\\cal P}(\\delta)$ is not necessarily a Gaussian field. Phase information is also important as a statistical diagnostic of non-linearity when the evolution of clustering in the non-linear regime which is generally intractable analytically. Indeed, , as we shall show, it is also closely related to the morphology of gravitational clustering and through this to the dynamical origin of structure.  The layout of this paper is as follows. In section 2 we present a visual demonstration of the link between Fourier phases and clustering morphology. Section 3 contains a brief discussion of some of the properties of phases and some pitfalls that must be avoided when using them as part of a clustering descriptor. In Section 4 we introduce a cross-correlation parameter $S$ and a rank-correlation $\\tau$ as quantitative measures of the agreement between two distributions. In Section 5 we display the results of the correlation tests between sample distributions and phase-based reconstructions, and between simulations evolving from different initial power spectra but the same initial phase set. A brief discussion of the results follows in Section 6. ",
        "conclusions": "We have shown the importance of Fourier phases on morphology of gravitational clustering via qualitative and quantitative demonstrations. It is interesting to remark upon the similarity of the results we have obtained here and those presented in Coles et al. (1993). The latter authors were interested in the ability of simple analytic methods to reproduce the clustering displayed by full N-body computations. They showed in particular that the Zel'dovich approximation (Zel'dovich 1970) could reproduce the full numerical results quite well, with cross-correlations similar to those we have found here. The Zel'dovich approximation works as well as it does in this respect because it places the caustic surfaces forming sheets and filaments near to the correct location in the N-body experiment. In other words, the Zel'dovich approximation has a high phase fidelity. It is less good at getting the amplitudes right. As we have shown, however, phases dominate the morphology.  With a new colour representation technique (Coles \\& Chiang 2000), phase information can be used to distinguish between non-Gaussianity induced by gravitational clustering and that by other mechanisms. For example, in the web page \\begin{verbatim} http://www.nottingham.ac.uk/~ppzpc/phases/cmb.html \\end{verbatim} the phase configuration of the non-Gaussian temperature fluctuations on the CMB sky induced by cosmic strings can be seen to be intrinsically different from that of hierarchical clustering developed in N-body simulations. A new algorithm based on the analysis of the phase distribution of Fourier components is recently devised to extract noise from point sources from the CMB signal(Naselsky et al. 2000), which is another example practice of the close link between Fourier phases and morphology.  Future large-scale galaxy redshift surveys and microwave sky maps will reveal much morphological information about large-scale structure in the Universe. Existing statistical technology, however, is still dominated by second-order methods that are blind to phase information. Our success in extracting this information will therefore depend on the development of statistical methods sufficiently sensitive to the key ingredient: the distribution of Fourier phases."
    },
    "0011/astro-ph0011217_arXiv.txt": {
        "abstract": "Can the submillimetre counts and background be produced by applying a locally derived extinction correction to the population of Lyman break galaxies? We investigate the submillimetre emission of two strongly lensed Lyman break galaxies ($\\sou{MS}{1512{+}36}$-cB58 and $\\sou{MS}{1358{+}62}$-G1) and find that the procedure that is used to predict the submillimetre emission of the Lyman break galaxy population overpredicts the observed $850\\mum$ fluxes by up to a factor of 14. This result calls for caution in applying local correlations to distant galaxies. It also shows that large extinction corrections on Lyman break galaxies should be viewed with skepticism. It is concluded that the Lyman break galaxies may contribute to the submillimetre background at the 25 to 50\\% level. The brighter submillimetre galaxies making up the rest of the background are either not detected in optical surveys, or if they are detected, their submillimetre emission cannot be reliably estimated from their rest-frame ultraviolet properties. ",
        "introduction": "\\label{sec.LBGs}  Measurements of the cosmic star formation density (SFD) based on surveys for Lyman break galaxies (LBGs) are affected by extinction, and attempts to correct for this effect lead to a substantial upwards revision of the SFD \\citep{Meureretal99}. Since the light absorbed in the ultraviolet (UV) is reradiated in the far-infrared (FIR), LBGs affected by extinction must emit FIR radiation, which will contribute to the submillimetre background.  However, the lack of plausible counterpart LBGs in deep submillimetre surveys with SCUBA at $850\\mum$ \\citep{Hughesetal98} has given rise to the view that the submillimetre galaxies form a separate population which is not represented in LBG samples in the first place, a point of view which has been the subject of considerable debate. We here discuss the relation between the LBGs and the submillimetre galaxies in the light of existing data on submillimetre source counts and background radiation, and recent measurements of submillimetre emission from LBGs, including lensed LBGs.  The expected submillimetre emission from the population of LBGs has recently been estimated based on the observed (if not entirely understood) correlation between the spectral index $\\beta$ in the UV (defined by the relation $f_\\lambda\\propto\\lambda^\\beta$) and the ratio of FIR to UV flux in a sample of local galaxies observed with the IUE satellite \\citep{Meureretal99}. Applying this relation to the LBG population, \\citet{AdelbergerSteidel00} found that the submillimetre counts and integrated background can be accounted for. As noted by these authors, these estimates are still uncertain, since the validity of the $\\beta$-FIR/UV correlation at high redshift has not been established, and important ingredients in the analysis such as the distribution of $\\beta$ values and its dependence on magnitude, and the luminosity function at faint magnitudes are poorly constrained. It is therefore necessary to verify this analysis by direct submillimetre observations of LBGs.  The predicted $850\\mum$ fluxes for most LBGs based on their UV properties \\citep{AdelbergerSteidel00} are $1\\un{mJy}$ or less, which is too faint for current instrumentation. However, some of the reddest LBGs are predicted to produce detectable $850\\mum$ emission and a sample of such LBGs can be selected based on their observed rest-frame UV color and magnitude. However, a deep search for $850\\mum$ emission from a sample of 8 LBGs predicted to have $S_{850}>1\\un{mJy}$ produced only one detection; it was concluded that the $\\beta$-FIR/UV correlation overpredicted the FIR emission by at least a factor of two \\citep{Chapmanetal00}. Recently, it has been argued, based on an analysis of structure in the faint $850\\mum$ emission of the Hubble Deep Field (HDF), that the LBGs contribute at least 25\\% but at most 50\\% of the submillimetre background \\citep{Peacocketal00}. These results indicate that the LBGs do contribute to the submillimetre background in a significant way, but that they do not produce a dominant contribution.  \\begin{table}[t] \\caption{Predicted (from UV color and magnitude) and observed submillimetre emission of lensed LBGs. Luminosities (not corrected for gravitational amplification) have been derived for $H_0=50\\kms\\pun{Mpc}{-1}$ and $q_0=0.5$. The other results do not depend on cosmology. Upper limits represent $3\\sigma$. \\label{tab.results}} \\vspace{0.4cm} \\begin{center} \\footnotesize \\begin{tabular}{|c|c|c|} \\hline {} &\\raisebox{0pt}[13pt][7pt]{$\\sou{MS}{1512{+}36}$-cB58} & \\raisebox{0pt}[13pt][7pt]{$\\sou{MS}{1358{+}62}$-G1}\\\\ \\hline & & \\\\[-20pt] amplification factor & $\\sim50$ & $\\sim10$ \\\\[-8pt] $\\beta$ (observed) & $-0.74\\pm0.1$ & $-1.63\\pm0.1$ \\\\[-8pt] $A_{1600}$ (predicted) & $3.0\\pm0.2$ & $1.2\\pm0.2$ \\\\[-8pt] $\\qu{S}{FIR}/\\qu{S}{1600}$ (predicted) & 18 & 2.4 \\\\[-8pt] $\\qu{L}{FIR}$ (predicted) & $3.9\\cdot10^{13}\\Lsun$ & $3.3\\cdot10^{12}\\Lsun$ \\\\[-8pt] $S_{850}$ (predicted) & $58\\un{mJy}$ & $5\\un{mJy}$ \\\\[-8pt] $S_{850}$ (observed)  & $4.2\\pm0.9\\un{mJy}$ & $<4\\un{mJy}$ \\\\[-8pt] $\\qu{L}{FIR}$ (derived) & $2.8\\cdot10^{12}\\Lsun$ & $<2.6\\cdot10^{12}\\Lsun$ \\\\[-8pt] $\\qu{S}{FIR}/\\qu{S}{1600}$ (derived) & 1.3 & $<1.9$ \\\\[-8pt] $A_{1600}$ (derived) & 0.8 & $<1.0$ \\\\[5pt] \\hline \\end{tabular} \\end{center} \\vspace{0.4cm} \\end{table} ",
        "conclusions": "These results demonstrate that attempts to produce the submillimetre counts and background based on an extinction correction applied to the LBG population \\citep{AdelbergerSteidel00} are fraught with considerable uncertainty, and fail in the case of the two lensed LBGs discussed here. While our sample is small, the results argue against the validity of the low redshift $\\beta$-FIR/UV correlation in the case of LBGs. Even if the local correlation were valid for high redshift galaxies as well, it still would not be able to produce the brighter submillimetre galaxies. These objects have luminosities putting them in the class of the ultraluminous infrared galaxies (ULIGs), which do not follow the $\\beta$-FIR/UV correlation, as shown by recent HST-STIS data of a sample of nearby ULIGs. Since these galaxies already account for $\\sim50\\%$ of the submillimetre background, it is not possible for the LBGs to produce a dominant fraction (let alone all) of the submillimetre background. A more likely situation is that the LBGs account for 25 to 50\\% of the submillimetre background as indicated by the faint structure in the HDF at $850\\mum$ \\citep{Peacocketal00}, but that the dominant part of the submillimetre background is made by a small number of ULIGs, which are either not detected in LBG surveys, or if they are detected, cannot be reliably corrected for extinction, in the same way that local ULIGs cannot be extinction corrected based on UV data.  In summary therefore, these results support the view that the brighter $850\\mum$ galaxies making at least 50\\% of the submillimetre background, form a population which cannot be reliably reproduced by extinction corrections applied to the LBG population."
    },
    "0011/astro-ph0011571_arXiv.txt": {
        "abstract": "\\noindent Low-mass disks orbiting a massive body can support ``slow'' normal modes, in which the eigenfrequency is much less than the orbital frequency. Slow modes are lopsided, i.e., the azimuthal wavenumber $m=1$. We investigate the properties of slow modes, using softened self-gravity as a simple model for collective effects in the disk. We employ both the WKB approximation and numerical solutions of the linear eigenvalue equation. We find that all slow modes are stable. Discrete slow modes can be divided into two types, which we label g-modes and p-modes. The g-modes involve long leading and long trailing waves, have properties determined by the self-gravity of the disk, and are only present in narrow rings or in disks where the precession rate is dominated by an external potential. In contrast, the properties of p-modes are determined by the interplay of self-gravity and other collective effects. P-modes involve both long and short waves, and in the WKB approximation appear in degenerate leading/trailing pairs. Disks support a finite number---sometimes zero---of discrete slow modes, and a continuum of singular modes. ",
        "introduction": "\\noindent Disks orbiting massive bodies are found in many astrophysical systems. These may be composed of gas (surrounding protostars, compact objects in binary star systems, or black holes in galactic nuclei), dust (a component of protoplanetary disks), solid bodies (planetary rings and planetesimal disks), or stars (galactic nuclei).  Normally the disk mass $M_d$ in these systems is much less than the central mass $M$, so that the gravitational potential is nearly Keplerian. A special feature of the Keplerian potential is that eccentric orbits do not precess. This feature arises from a 1:1 resonance between the radial and azimuthal frequencies, and implies that the evolution of the eccentricity distribution in an isolated Keplerian disk is determined by the collective effects in the disk (self-gravity, pressure, collisions, viscosity, etc.), no matter how small these may be. Thus we expect that Keplerian disks may be capable of supporting a set of normal modes with eigenfrequency $\\omega$ much less than the Keplerian orbital frequency $\\Omega(r)=(GM/r^3)^{1/2}$, and properties determined by weak collective effects in the disk. In the linear limit these modes have azimuthal wavenumber $m=1$. We shall call them ``slow'' modes.  The goal of this paper is to investigate slow modes in Keplerian disks.  In this Section we provide an introduction and orientation based on the WKB approximation. We shall find that in many cases slow modes have scales comparable to the size of the disk, so that the WKB analysis must be supplemented by numerical calculations of large-scale modes. We soften the self-gravity of the disk to mimic the effects of pressure in fluid disks and velocity dispersion in collisionless disks, and derive the integral eigenvalue equation that describes large-scale slow modes in \\S II. We discuss the numerical solution of this integral equation for several disk models in \\S III. Section IV contains a discussion of earlier related work, and \\S V contains conclusions.  \\subsection{The precession rate}  \\noindent We employ cylindrical coordinates $(r,\\phi,z)$ with origin at the central mass $M$, and approximate the unperturbed disk as razor-thin, with surface density $\\Sigma_d(r)$ in the $z=0$ plane. We assume that the disk mass is small, $M_d/M\\sim\\Sigma_dr^2/M\\sim \\epsilon\\ll1$. The unperturbed disk material orbits in the axisymmetric potential \\be \\Phi(r)=-{GM\\over r}+\\Phi_d(r)+\\Phi_e(r), \\ee where $\\Phi_d(r)$ is the potential arising from the self-gravity of the disk and $\\Phi_e(r)$ is the non-Keplerian potential from any external source, and both are assumed to be O$(\\epsilon)$.  For nearly circular orbits, the azimuthal and radial frequencies $\\Omega>0$ and $\\kappa>0$ are given by (e.g. \\nocite{bt87}Binney and Tremaine 1987) \\begin{eqnarray} \\Omega^2(r) & = & {GM\\over r^3}+{1\\over r}{d\\over dr}(\\Phi_d+\\Phi_e), \\nonumber \\\\ \\kappa^2(r) & = & {GM\\over r^3}+{d^2\\over dr^2}(\\Phi_d+\\Phi_e)+ {3\\over r}{d\\over dr}(\\Phi_d+\\Phi_e). \\label{eq:kapdef} \\end{eqnarray} The line of apsides of an eccentric test-particle orbit subjected only to gravity precesses at a rate \\begin{eqnarray} \\dot\\varpi & = & \\Omega-\\kappa \\nonumber \\\\ & = &-{1\\over 2\\Omega(r)}\\left({2\\over r}{d\\over dr} +{d^2\\over dr^2}\\right)(\\Phi_d+\\Phi_e)+\\hbox{O}(\\epsilon^2), \\label{eq:precdef} \\end{eqnarray} where from now on $\\Omega(r)$ is assigned the Keplerian value, $(GM/r^3)^{1/2}$.  \\subsection{The Kuzmin disk}  \\label{sec:kuzmin}  \\noindent Several of our examples will be based on the Kuzmin disk, which has surface density and potential (e.g. \\nocite{bt87}Binney and Tremaine 1987) \\be \\Sigma_d(r)={aM_d\\over2\\pi(r^2+a^2)^{3/2}}, \\qquad \\Phi_d(r)=-{GM_d\\over (r^2+a^2)^{1/2}}, \\label{eq:kuzmin} \\ee where $M_d$ is the disk mass and $a$ is the disk scale length. Using equation (\\ref{eq:precdef}), the precession rate due to the Kuzmin disk is \\be \\dot\\varpi_d=-{3GM_da^2\\over 2\\Omega(r)(r^2+a^2)^{5/2}}, \\ee for $M_d\\ll M$.  The precession rate is negative at all radii, approaches zero as $r^{3/2}$ for small radii and as $r^{-7/2}$ at large radii, and has a single extremum at $r_0=(\\frac{3}{7})^{1/2}a=0.6547a$ where $\\dot\\varpi(r_0)=-0.3257(M_d/M)(GM/a^3)^{1/2}$.  \\subsection{The WKB approximation}  \\label{sec:wkb}  \\noindent We can write the linearized surface-density response of the disk as a superposition of terms of the form $\\Sigma(r,\\phi,t)=A(r)\\exp\\left\\{i[\\int^rk(r)dr+m\\phi-\\omega t]\\right\\}$, where $m$ is the azimuthal wavenumber, $k(r)$ is the radial wavenumber, and $\\omega$ is the frequency. In the WKB or tight-winding approximation where $|kr|\\gg 1$, the dispersion relation between $k(r)$ and $\\omega$ depends only on the local properties of the disk, and can be derived analytically for both fluid and collisionless disks.  For a barotropic fluid disk with sound speed $c(r)$, the WKB dispersion relation reads (Safronov 1960; Binney and Tremaine 1987)\\nocite{saf60,bt87} \\be (\\omega-m\\Omega)^2=\\kappa^2-2\\pi G\\Sigma_d|k|+k^2c^2. \\label{eq:disprel} \\ee This relation implies that the disk is locally stable to axisymmetric ($m=0$) disturbances if and only if \\be Q\\equiv {c\\kappa\\over \\pi G\\Sigma_d} > 1. \\label{eq:qgas} \\ee For thin, low-mass disks we have $c\\ll\\Omega r$ and $\\Sigma_d\\ll M/r^2$, so that in general the dispersion relation (\\ref{eq:disprel}) can only be satisfied if $|kr|\\gg 1$, which justifies the use of the WKB approximation. However, the special case of $m=1$ disturbances in a nearly Keplerian disk is different. Using equation (\\ref{eq:precdef}) the dispersion relation can be rewritten in the form \\be \\omega=\\dot\\varpi +{\\pi G\\Sigma_d|k|\\over\\Omega} - {k^2c^2\\over 2\\Omega} +{1\\over\\Omega}\\hbox{O}(\\dot\\varpi^2,\\omega^2). \\label{eq:dispone} \\ee Since we are interested in slow modes in nearly Keplerian disks, the terms of order $\\dot\\varpi^2$, $\\omega^2$ can be dropped. At a given radius, the maximum value of $\\omega$ is attained at wavenumber $|k|\\equiv k_0=\\pi G\\Sigma_d/c^2$.  Let us neglect the effects of pressure for the moment, by dropping the term proportional to $c^2$ in equation (\\ref{eq:dispone}).  Thus the WKB dispersion relation for slow waves dominated by self-gravity becomes \\be \\omega=\\dot\\varpi +{\\pi G\\Sigma_d|k|\\over\\Omega}. \\label{eq:wkb} \\ee Since $\\dot\\varpi/\\Omega=\\hbox{O}(\\Sigma_dr^2/M)$ for isolated disks, in general the two terms on the right side of (\\ref{eq:dispone}) are only comparable if $|kr|\\sim 1$; thus slow gravity-dominated waves must be large-scale (i.e. the WKB approximation is {\\em not} satisfied) and have eigenfrequency $\\omega=\\hbox{O}(\\dot\\varpi)$. The neglect of the pressure term is valid for such waves so long as $G\\Sigma_dr\\gg c^2$, or ${\\cal M}\\gg Q$, where ${\\cal M}(r)\\equiv\\Omega r/c$ is the Mach number.  Next consider a collisionless disk with a Schwarzschild distribution function having radial velocity dispersion $c_R(r)$. The WKB dispersion relation reads (Kalnajs 1965; Lin and Shu 1966; Binney and Tremaine 1987)\\nocite{kal65,ls66,bt87} \\be (\\omega-m\\Omega)^2=\\kappa^2-2\\pi G\\Sigma_d|k|{\\cal F}\\left({\\omega-m\\Omega\\over\\kappa},{k^2c_R^2\\over\\kappa^2}\\right), \\label{eq:disprels} \\ee where ${\\cal F}(s,\\chi)$ is the reduction factor. Explicit formulae for the reduction factor are given in the references above. The disk is locally stable to axisymmetric disturbances if and only if (Toomre 1964)\\nocite{too64} \\be Q\\equiv {c_R\\kappa\\over 3.36 G\\Sigma_d} > 1, \\ee which closely resembles equation (\\ref{eq:qgas}).  For $m=1$ disturbances in a nearly Keplerian disk, the dispersion relation (\\ref{eq:disprels}) simplifies to \\be \\omega=\\dot\\varpi +{\\pi G\\Sigma_d|k|\\over\\Omega}{\\cal F}_K \\left({k^2c_R^2\\over\\Omega^2}\\right), \\label{eq:dispxxx} \\ee in which \\be {\\cal F}_K(\\chi)={2\\over\\chi}e^{-\\chi}I_1(\\chi), \\ee and $I_1(\\chi)$ is a modified Bessel function. At a given radius, the maximum value of $\\omega$ is attained at wavenumber $|k|\\equiv k_0=0.7643\\,\\Omega/c_R$. Since both terms on the right side of (\\ref{eq:dispxxx}) scale as the disk surface density $\\Sigma_d$, for a fixed velocity-dispersion profile the shape of a mode is independent of the disk mass and the eigenfrequency is proportional to the disk mass.  For the moment assume $\\chi\\ll1$; then the reduction factor ${\\cal F}(\\chi)=1$ and the dispersion relation (\\ref{eq:dispxxx}) reduces to (\\ref{eq:wkb}): we conclude, not surprisingly, that the WKB dispersion relation for slow gravity waves is the same for fluid and collisionless disks. The neglect of the reduction factor is valid for such waves so long as $\\chi\\ll1$ when $|kr|\\sim 1$, which in turn requires that the Mach number ${\\cal M}_R\\equiv \\Omega r/c_R\\gg 1$.  Many of the properties of waves in collisionless disks can be mimicked by softening the gravitational potential, using $-Gm/(d^2+b^2)^{1/2}$ instead of $-Gm/d$ for the potential due to a mass $m$ at distance $d$. Here the parameter $b$ is the softening length. This modification yields the WKB dispersion relation (Miller 1971; Erickson 1974; Toomre 1977)\\nocite{mil71,eri74,too77} \\be \\omega=\\dot\\varpi(r) +{\\pi G\\Sigma_d(r)\\over\\Omega(r)}|k|\\exp(-|k|b). \\label{eq:dispsoft} \\ee At a given radius, the maximum value of $\\omega$ is attained at wavenumber $|k|\\equiv k_0=1/b$. As in the case of a collisionless disk, both terms on the right side scale as the disk surface density $\\Sigma_d$, so the eigenfrequency is proportional to the disk mass. In the long-wavelength limit, $|k|b\\ll1$, the dispersion relation simplifies once again to equation (\\ref{eq:wkb}).  By analogy with the theory of WKB waves in galactic disks, we call waves trailing if $k>0$ and leading if $k<0$. Long waves have $|k|<k_0$ and short waves have $|k|>k_0$.  We conclude from these approximate arguments that a low-mass, thin disk (by which we mean $\\Sigma_dr^2/M\\sim\\epsilon\\ll1$, and ${\\cal M}\\gg Q$ for a fluid disk, or ${\\cal M}\\gg 1$ for a collisionless disk) could support slow modes with the following properties: (i) the azimuthal wavenumber $m=1$; (ii) the dominant collective effect is self-gravity; (iii) $|kr|=\\hbox{O}(1)$, that is, the wavelength is comparable to the disk radius, the properties of the mode are determined globally rather than locally, and the mode cannot be described accurately with the WKB approximation; (iv) the properties of the mode depend on the disk self-gravity but not on other collective effects; (v) the eigenfrequency is of order $\\epsilon\\Omega$. Many of these conclusions were derived already by \\nocite{lg99}Lee and Goodman (1999), in a paper focused mainly on nonlinear waves in nearly Keplerian disks.  \\begin{figure} \\plotone{Tremaine.fig1.ps} \\caption{Contours of constant frequency $\\omega$ for WKB slow waves in a Kuzmin disk (\\S \\ref{sec:kuzmin}), as a function of wavenumber $kr$ and radius $r$. The dispersion relations shown are for a disk with softened gravity (eq. \\ref{eq:dispsoft}), a fluid disk (eq. \\ref{eq:dispone}), a collisionless disk (eq. \\ref{eq:dispxxx}), and a gravity-dominated disk (eq. \\ref{eq:wkb}). The disk scalelength $a=1$, the softening length $b=0.2r$, the Mach number ${\\cal M}=\\Omega /cr$ or ${\\cal M}_R=\\Omega r/c_R$ is the same for the fluid and the collisionless disk and equal to 5 at all radii.  Both the gravitational constant $G$ and the central mass $M$ are unity; the disk mass $M_d=1$ and $\\omega$ scales as $M_d$ so long as $b$, ${\\cal M}_R$ or $M_d{\\cal M}^2$ remain invariant. The extrema at $kr=0$ are minima. The contours in the right-hand panels are not uniformly spaced.} \\label{fig:wkb} \\end{figure}  Despite conclusion (iii), the WKB approximation provides an invaluable guide to the properties of slow modes.  It is useful to plot contours of constant $\\omega$ in the $(k,r)$ plane. Figure \\ref{fig:wkb} shows such contour plots for the Kuzmin disk discussed in \\S \\ref{sec:kuzmin} and four dispersion relations for slow waves: a disk with softened gravity (eq.~\\ref{eq:dispsoft}), a fluid disk (eq.~\\ref{eq:dispone}), a collisionless disk (eq.~\\ref{eq:dispxxx}), and a gravity-dominated disk (eq.~\\ref{eq:wkb}). Local maxima in the contour plots occur at $k=\\pm k_0$, and a local minimum at $k=0$.  In these diagrams, a WKB wavepacket travels along a contour of constant $\\omega$, at a radial speed given by the group velocity ${d\\omega/ dk}$ \\nocite{too69}(Toomre 1969). Long trailing and short leading waves propagate outwards, while long leading and short trailing waves propagate in.  In the WKB approximation a mode corresponds to a closed contour in Figure \\ref{fig:wkb} on which the total phase change $\\oint k(r)dr$ over one circuit is an integer multiple of $2\\pi$, including possible phase shifts at the turning points. The models in the Figure exhibit two types of closed contour:  \\begin{enumerate}  \\item On contours such as C in the upper-left panel, an outward-traveling long trailing wave refracts into an inward-traveling short trailing wave, which in turn refracts back into a long wave. Contours such as B are similar, except that they involve inward-traveling long leading waves and outward-traveling short leading waves. WKB modes of types B and C occur in degenerate pairs. Modes of this kind are not present in the gravity-dominated disk (lower right panel), since this disk supports only long waves. In the Kuzmin disk these modes have $\\omega>0$, and the modes with the smallest phase change around the contour have the highest frequency and wavenumbers localized near $\\pm k_0$. We call these {\\em p-modes}, since they require the presence of pressure, velocity dispersion, or softening.  \\item In contrast, closed contours such as A involve only long waves, and involve reflection of an outward-traveling long trailing wave into an inward-traveling long leading wave, which then reflects back to long trailing. These modes are present in all four panels; their properties are similar in all panels since they depend only on the surface-density distribution in the disk and not other collective effects. In the Kuzmin disk these modes have $\\omega<0$, and the modes with the fewest nodes have the most negative values of $\\omega$.  The WKB approximation is highly suspect for waves of this kind since $|kr|\\lesssim 1$, so numerical mode calculations are required. We call these {\\em g-modes}, since self-gravity is the only collective effect involved.  \\end{enumerate}  Figure \\ref{fig:wkb} suggests that the WKB dispersion relation for collisionless Keplerian disks can be reproduced quite well by disks with softened gravity \\nocite{eri74,too77}(Erickson 1974; Toomre 1977); and at least the qualitative features of the dispersion relation in fluid disks are similar as well. Therefore, in our numerical mode calculations we shall soften the self-gravity of the disk but neglect other collective effects. ",
        "conclusions": "\\noindent Disks orbiting massive bodies support slow modes; these are linear $m=1$ normal modes whose eigenfrequency is proportional to the strength of collective effects in the disk, rather than to its characteristic radial or azimuthal frequencies. The most important collective effects are likely to be the self-gravity of the disk, and the velocity dispersion in collisionless disks or the pressure in fluid disks. The first two of these---and to a lesser extent the third---can be approximated by softened self-gravity (Fig. \\ref{fig:wkb}). Slow modes are important because these are likely to be the only large-scale or long-wavelength modes, and hence can dominate the overall appearance of the disk. Moreover, slow modes are relatively immune to damping by viscosity, Landau damping, collisions, or other dissipative mechanisms.  Slow modes in which the disk self-gravity is the dominant collective effect are described by the same equations used in the classic Laplace-Lagrange secular perturbation theory for planetary systems, except that (i) the number of planets $N$ goes to infinity and the mass per planet $m\\propto N^{-1}$; (ii) the Laplace-Lagrange equations are valid in the limit where the planets' radial excursions are much less than their radial separation, whereas in continuous disks the radial separation of the mass elements is smaller than the radial excursions. The simplest way to modify Laplace-Lagrange theory to account for this difference is to soften gravity.  We have derived the eigenvalue equation (\\ref{eq:intsymma}) that describes slow modes in a nearly Keplerian disk with softened gravity. We have solved this equation numerically to investigate the properties of modes in the Kuzmin disk (eq. \\ref{eq:kuzmin}), the Gaussian ring (eq. \\ref{eq:gauss}), the Jacobs-Sellwood ring (eq. \\ref{eq:sell}), and a variety of other disks not reported here. Many of our results can be interpreted in the framework of the WKB approximation (\\S \\ref{sec:wkb}). Our conclusions can be summarized as follows:  \\begin{enumerate}  \\item All slow modes have real frequency, and thus are stable.  \\item There are two main types of discrete slow mode: (i) g-modes involve long leading and long trailing waves. Their properties are insensitive to the softening parameter $\\beta$ once it is sufficiently small. In general these are retrograde (eigenfrequency $\\omega<0$). The fundamental g-mode has zero nodes, the lowest eigenfrequency, and a scale comparable to the disk radius. (ii) The properties of p-modes depend on the softening parameter; in particular their characteristic wavenumber is $k\\sim\\pm\\beta/r$. In the WKB approximation that is valid for $\\beta\\ll1$, p-modes come in degenerate pairs, one composed of short and long trailing waves and the other of short and long leading waves. In general p-modes are prograde ($\\omega>0$), and the fundamental p-mode has the largest $\\omega$.  \\item Narrow rings have a trivial slow g-mode, $e_K(r)=$constant. We have examined a wide variety of broad rings and disks, but have found none that support g-modes when they are isolated. In the presence of an external non-Keplerian potential that supplies a fraction $1-\\lambda$ of the total precession, a variety of disks appear to support g-modes if and only if $\\lambda\\lesssim 0.5$--0.6.  \\item The p-modes are modulated at the characteristic spatial frequency $k_0=\\beta/r$. Thus long-wavelength p-modes occur only in disks with substantial softening (cf. Figures \\ref{fig:shapp}, \\ref{fig:gauss}, and \\ref{fig:sellmode}).  \\item In general disks support a finite number of discrete normal modes and a continuum of singular modes. In some cases, such as the Jacobs-Sellwood ring with $\\beta>0.229$, there are no discrete normal modes.  \\end{enumerate}  Of course, these models have many shortcomings: they are linearized; they use softened gravity as a crude proxy for velocity dispersion or pressure; and the softening law (eq. \\ref{eq:ooppsoft}) was chosen for mathematical convenience and does not correspond to a simple interparticle force law. The models also do not address the important questions of how rapidly slow modes damp (although the N-body simulations of Jacobs and Sellwood 1999\\nocite{js99} suggest they can be very long-lived), or how large the disk-mass ratio $M_d/M$ can be in disks that support discrete slow modes. To decide these issues there is no substitute for careful $N$-body simulations; however, our results provide a guide for interpreting these simulations. This paper also does not address the question of how slow modes are excited and whether long-lived slow modes actually play a significant role in the structure and evolution of a variety of astrophysical disks."
    },
    "0011/astro-ph0011092_arXiv.txt": {
        "abstract": "We present the luminosity function and color-redshift relation of a magnitude-limited sample of 145 mostly red field E/S0 galaxies at $z \\lesssim 1$ from the DEEP Groth Strip Survey (GSS). Using nearby galaxy images as a training set, we develop a quantitative method  to classify E/S0 galaxies based on smoothness, symmetry, and bulge-to-total light ratio. Using this method, we identify 145 E/S0s at $16.5 < I < 22$ within the GSS, for which 44 spectroscopic redshifts ($z_{spec}$) are available. Most of the galaxies with spectroscopic redshifts  (86\\%) form a {\\it red envelope} in the redshift-color diagram, consistent with predictions of spectral synthesis models in which the dominant stellar population is formed at redshifts $z \\gtrsim 1.5$. We use the tight correlation between $V-I$ and $z_{spec}$ for this red subset to estimate redshifts of the remaining E/S0s to an accuracy of $\\sim$10\\%, with the exception of a small number (16\\%) of blue interlopers at low redshift that are quantitatively classified as E/S0s but are not contained within the red envelope. Constructing a luminosity function of the full sample of 145 E/S0s, we find that there is about 1.1--1.9 magnitude brightening in rest-frame $B$ band luminosity back to  $z \\simeq 0.8$ from $z=0$, consistent with other studies. Together with the red colors, this brightening is consistent with models in which the bulk of stars in red field E/S0s formed before $z_{for} \\gtrsim 1.5$ and have been evolving rather quiescently with few large starbursts since then. Evolution in the number density of field E/S0 galaxies is harder to measure, and uncertainties in the raw counts and their ratio to local samples might amount to as much as a factor of two. Within that uncertainty, the number density of red E/S0s to $z \\simeq 0.8$ seems relatively static, being comparable to or perhaps moderately less than that of local E/S0s depending on the assumed cosmology. A doubling of E/S0 number density since $z = 1$ can be ruled out with high confidence (97\\%) if $\\Omega_{m}=1$. Taken together, our results are consistent with the hypothesis that the majority of luminous field E/S0s were already in place by $z \\sim 1$, that the bulk of their stars were already fairly old, and that their number density has not changed by large amounts since then. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0011/astro-ph0011437_arXiv.txt": {
        "abstract": "We report results from our mini-survey for molecular hydrogen in eight high-redshift damped Lyman-$\\alpha$ (DLA) systems using the ESO Ultra-violet and Visible Spectrograph on the VLT. In addition, we investigate two systems using ESO public data. We include in the sample the only system where \\h2 was previously detected and studied at high-spectral resolution. Altogether our sample consists of eleven absorbers with 1.85~$<$~$z_{\\rm abs}$~$<$~3.4.\\par\\noindent We confirm the presence of \\h2 in the $z_{\\rm abs}$ = 2.3377, metal-poor ([Si/H] = $-$1.20), system toward PKS 1232+082. The derived molecular fraction, $f$~= 2$N$(H$_2$)/(2$N$(H$_2$)+$N$(H~{\\sc i})) = 4$\\times$10$^{-4}$, is two orders of magnitude less than what has been claimed previously from low-resolution data. The physical conditions within the cloud can be constrained directly from observation. The kinetic temperature and particle density are in the ranges, respectively, 100~$<$~$T$~$<$~300~K and 30~$<$~$n_{\\sc H}$~$<$~50~cm$^{-3}$. In addition, UV pumping is of the same order of magnitude than in our Galaxy.\\par\\noindent The upper limits on the molecular fraction derived in nine of the systems are in the range 1.2$\\times$10$^{-7}$$-$1.6$\\times$10$^{-5}$. There is no evidence in this sample for any correlation between \\h2 abundance and relative heavy element depletion into dust grains. This should be investigated using a larger sample however. The molecular abundance in a few DLA systems (and in particular in the two systems where \\h2 is detected) is consistent with what is seen in the Magellanic clouds. But most of the DLA measurements are well below these values. This is probably partly due to small amounts of dust and/or high UV flux. We argue however that the lack of molecules is a direct consequence of high kinetic temperature ($T$~$>$~3000~K) implying a low formation rate of \\h2 onto dust grains. Therefore, most of the DLA systems arise in warm and diffuse neutral gas. ",
        "introduction": "The Ultra-violet and Visible Echelle Spectrograph (D'Odorico et al. 2000) mounted on the ESO Kueyen 8.2~m telescope at the Paranal observatory has been used on April 5 to 8 and 29, 2000 to obtain high-spectral resolution spectra of eight previously known DLA systems at $z_{\\rm abs}$~$>$~1.8. Settings have been adjusted in both arms to cover the \\h2 Werner and Lyman series absorption range together with several important metal absorption lines. The slit width was 1~arcsec (the seeing $FWHM$ was most of the time better than 0.8~arcsec) and the CCDs were binned 2$\\times$2 resulting in a resolution of $\\sim$45000. The exposure time $\\sim$2~hours was split into two exposures. The data was reduced using the UVES pipeline, a set of procedures implemented in a dedicated context of MIDAS, the ESO data reduction package. The main characteristics of the pipeline is to perform a precise inter-order background subtraction for science frames and master flat-fields, and an optimal extraction of the object signal rejecting cosmic ray impacts and subtracting the sky at the same time. Wavelengths were corrected to vacuum-heliocentric values and individual 1D spectra were combined together. The resulting S/N ratio per pixel is of the order of 10 at $\\sim$3500~\\AA~ and  20 at $\\sim$6000~\\AA. \\vskip -0.3cm ",
        "conclusions": ""
    },
    "0011/hep-th0011014_arXiv.txt": {
        "abstract": "\\noindent After the separation of the center-of-mass motion, a new privileged class of canonical Darboux bases is proposed for the non-relativistic N-body problem by exploiting a geometrical and group theoretical approach to the definition of {\\it body frame} for deformable bodies. This basis is adapted to the rotation group SO(3), whose canonical realization is associated with a symmetry Hamiltonian {\\it left action}. The analysis of the SO(3) coadjoint orbits contained in the N-body phase space implies the existence of a {\\it spin frame} for the N-body system. Then, the existence of appropriate non-symmetry Hamiltonian {\\it right actions} for non-rigid systems leads to the construction of a N-dependent discrete number of {\\it dynamical body frames} for the N-body system, hence to the associated notions of {\\it dynamical} and {\\it measurable} orientation and shape variables, angular velocity, rotational and vibrational configurations. For N=3 the dynamical body frame turns out to be unique and our approach reproduces the {\\it xxzz gauge} of the gauge theory associated with the {\\it orientation-shape} SO(3) principal bundle approach of Littlejohn and Reinsch. For $N \\geq 4$ our description is different, since the dynamical body frames turn out to be {\\it momentum dependent}. The resulting Darboux bases for $N\\geq 4$ are connected to the coupling of the {\\it spins} of particle clusters rather than the coupling of the {\\it centers of mass} (based on Jacobi relative normal coordinates). One  of the advantages of the spin coupling is that, unlike the center-of-mass coupling, it admits a relativistic generalization.    \\vskip 1truecm  \\today    \\vskip 1truecm ",
        "introduction": "This paper deals with the construction of a specialized system of coordinates for the non-relativistic N-body problem which could be instrumental to nuclear, atomic and molecular physics, as well as to celestial mechanics. In particular, we shall exploit the technique of the canonical realizations of Lie symmetry groups \\cite{pauri2,pauri1,pauri3,pauri4,pauri5} within the framework of the non-relativistic version of the {\\it rest-frame Wigner covariant instant form} of dynamics \\cite{lu,india} to the effect of obtaining coordinates adapted (locally in general) to the SO(3) group. In most of the paper we consider only free particles, since the mutual interactions are irrelevant to the definition of the kinematics in the non-relativistic case.  Isolated systems of  N particles possess 3N degrees of freedom in configuration space and 6N in phase space. The Abelian nature of the overall translational invariance, with its associated three commuting Noether constants of the motion, allows for the decoupling and, therefore, for the elimination of either three configurational variables or three pairs of canonical variables, respectively (separation of the center-of-mass motion). In this way one is left with either 3N-3 relative coordinates ${\\vec \\rho}_a$ or 6N-6 relative phase space variables ${\\vec \\rho}_a$, ${\\vec p}_a$, $a=1,.., N-1$ and the center-of-mass angular momentum or spin is $\\vec S= \\sum_{a=1}^{N-1}{\\vec \\rho}_a \\times {\\vec p}_a$. In the non-relativistic theory most of the calculations employ the sets of $3N-3$ {\\it Jacobi normal relative coordinates} ${\\vec s}_a$ which diagonalize the quadratic form associated with the  relative kinetic energy [the spin becomes $\\vec S=\\sum_{a=1}^{N-1}{\\vec s}_a \\times {\\vec \\pi}_{sa}$, with ${\\vec \\pi}_{sa}$  momenta conjugated to the ${\\vec s}_a$'s]. Each set of relative Jacobi normal coordinates ${\\vec s}_a$, $a=1,..,N-1$, is associated with a {\\it different clustering} of the N particles, corresponding to  the centers of mass of the various subclusters. In special relativity Jacobi normal coordinates do not exist, as it will be shown in Ref.\\cite{iten1}, and a different strategy must be used.     On the other hand, the non-Abelian nature of the overall rotational invariance entails the impossibility of an analogous intrinsic separation of {\\it rotational} (or {\\it orientational}) configurational variables from others which could be called {\\it shape} or {\\it vibrational}. As a matter of fact, this is one of the main concerns of molecular physics and of advanced mechanics of deformable bodies. Recently, a new approach inspired by the geometrical techniques of fiber bundles has been proposed in these fields of research: a self-contained and comprehensive exposition  of this viewpoint and a rich bibliography can be found in Ref.\\cite{little}.  In the theory of deformable bodies one looses any intrinsic notion of {\\it body frame}, which is a fundamental tool for the description of rigid bodies and  their associated Euler equations. A priori, for a given configuration of  a non-relativistic continuous body, and in particular for a N-body system, any barycentric orthogonal frame could be named {\\it body frame} of the system with its associated notion of {\\it vibrations}.  This state of affairs suggested \\cite{little} to replace the kinematically accessible region of the non-singular configurations \\footnote{ See Refs.\\cite{little,little1} for a discussion of the singular (collinear and N-body collision) configurations.  Applying the SO(3) operations to any given configuration of the 3N-3 relative variables (a point in the relative configuration space) gives rise to 3 possibilities only: i) for {\\it generic configurations} the orbit containing all the rotated copies of the configuration is a 3-dimensional manifold [diffeomorphic to the group manifold of SO(3)]; ii) for {\\it collinear configurations} the orbit is diffeomorphic to the 2-sphere $S^2$; iii) for the {\\it N-body collision configuration} (in which all the particles coincide at a single point in space) the orbit is a point.} in the (3N-3)-dimensional relative configuration space by a SO(3) principal fiber bundle over a (3N-6)-dimensional base manifold, called {\\it shape} space \\footnote{It is the open set of all the orbits of {\\it generic non-singular configurations}.}. The SO(3) fiber on each shape configuration carries the {\\it orientational} variables (e.g. the usual Euler angles) referred to the chosen {\\it body frame}.  A local cross section of the principal fiber bundle selects just one orientation of a generic N-body configuration in each fiber (SO(3) orbit) and this is equivalent to a {\\it gauge convention}, namely to a possible definition of a {\\it body frame} ({\\it reference orientation}), to be adopted after a preliminary choice of the {\\it shape} variables. It turns out that this principal bundle is trivial only for N=3, so that in this case only  global cross sections exist, and in particular the identity cross section may be identified with the {\\it space frame}. In this case any global cross section is a copy of the 3-body shape space and its coordinatization gives a description of the {\\it internal vibrational} motions associated with the chosen gauge convention for the reference orientation. For $N\\geq 4$, however, global cross sections do not exist \\footnote{This is due to the topological complexity of the shape space generated by the singular configurations \\cite{little1}, which are dispersed among the generic configurations for $N \\geq 4$.} and the definition of the reference orientation ({\\it body frame}) can be given only locally. This means that the  shape space  cannot be identified with a (3N-6)-dimensional submanifold of the (3N-3)-dimensional relative configuration space. The {\\it gauge convention} about the reference orientation and the consequent individuation of the internal vibrational degrees of freedom requires the choice of a connection on the SO(3) principal bundle (i.e. a concept of horizontality) and this leads in turn to the introduction of a SO(3) gauge potential on the base manifold. Obviously, physical quantities like the rotational or vibrational kinetic energies and, in general, any observable feature of the system must be gauge invariant, namely independent of the chosen convention. Note that both the space frame and the body frame components of the angular velocity are gauge quantities in the {\\it orientation-shape bundle} approach and their definition depends upon the gauge convention.  While a natural gauge invariant concept of {\\it purely rotational} N-body configurations exists (when the N-body velocity vector field is {\\it vertical}, i.e. when the shape velocities vanish), a notion of {\\it horizontal} or {\\it purely vibrational} configuration requires the introduction a connection $\\Gamma$ on the SO(3) principal bundle. A gauge fixing is needed in addition in order to select a particular $\\Gamma$-horizontal cross section and the correlated  gauge potential on the shape space. See Ref.\\cite{little} for a review of the gauge fixings used in molecular physics' literature  and, in particular, for the virtues of a special connection C corresponding to the shape configurations with vanishing center-of-mass angular momentum $\\vec S$ \\footnote{The C-horizontal cross sections are orthogonal to the fibers with respect to the Riemannian metric dictated by the kinetic energy.}.  This {\\it orientation-shape} approach replaces the usual Euler kinematics of rigid bodies and entails in general a coupling between the internal {\\it shape} variables and some of the {\\it orientational} degrees of freedom. In Ref.\\cite{little} it is interestingly shown that the non-triviality of the SO(3) principal bundle, when extended to continuous deformable bodies, is at the heart of the explanation of problems like the {\\it falling cat} and the {\\it diver}. A characteristic role of SO(3) gauge potentials in this case is to {\\it generate rotations by changing the shape}.  In Ref.\\cite{little}  the Hamiltonian formulation of this framework is also given, but no explicit procedure for the construction of a canonical Darboux basis for the orientational and shape variables is worked out. See Refs.\\cite{little,little1} for the existing sets of shape variables for $N=3,4$ and for the determination of their physical domain.  Independently of this SO(3) principal bundle framework and having in mind the relativistic N-body problem where only Hamiltonian methods are available, we have been induced to search for a {\\it constructive procedure} for building canonical Darboux bases in the (6N-6)-dimensional relative phase space, suited to the non-Abelian canonical reduction of the overall rotational symmetry. Our procedure surfaced from the following independent pieces of information:      A) In recent years a systematic study of relativistic kinematics of the N-body problem, in the framework of the {\\it rest-frame Wigner covariant instant form of dynamics} has been developed in Ref.\\cite{lus} and then applied to the isolated system composed  by N scalar charged particles plus the electromagnetic field \\cite{albad,crater}.  These papers contain the construction of a special class of canonical transformations, of the Shanmugadhasan type \\cite{sha,lu}. These transformations are simultaneously adapted to: i) the Dirac first class constraints appearing in the  Hamiltonian formulation of relativistic models; ii) the timelike Poincar\\'e orbits associated with most of their configurations. In the Darboux bases one of the final canonical variables is the square root of the Poincar\\'e invariant $P^2$ ($P_{\\mu}$  is the conserved timelike four-momentum of the isolated system). Subsequently, by using the constructive theory of the canonical realizations of Lie groups \\cite{pauri2,pauri1,pauri3,pauri4,pauri5} a new family of canonical transformations was introduced in Ref.\\cite{lucenti}. This latter leads to the definition of the so-called {\\it canonical spin bases}, in which also the Pauli-Lubanski Poincar\\'e  invariant $W^2=-P^2 {\\vec S}^2_T$ for timelike Poincar\\'e orbits \\footnote{For the configurations of the isolated system having a rest-frame Thomas canonical spin ${\\vec S}_T$ different from zero.} becomes one of the final canonical variables. The construction of the spin bases exploits the clustering of spins rather than the Jacobi clustering of centers of mass.  In spite of its genesis in a relativistic context, the technique used in the determination  of the spin bases, related to a {\\it typical form} \\cite{pauri2} of the canonical realizations of the E(3) group, can be easily adapted to the non-relativistic case, where $W^2$ is replaced by the invariant ${\\vec S}^2$ of the extended Galilei group.     B) These results provide the starting point for the construction of a canonical Darboux basis adapted to the non-Abelian SO(3) symmetry. The three non-Abelian Noether constants of motion $\\vec S$ are arranged in these canonical Darboux bases as an array containing the canonical pair $S^3$, $\\beta = tg^{-1}{{S^2}\\over {S^1}}$ and the unpaired variable $S=|\\vec S|$ \\footnote{In this context the configurations with $\\vec S =0$ are singular and have to be treated separately.} ({\\it scheme A} of the canonical realization of SO(3) \\cite{pauri1}). The angle canonically conjugated to $S$, say $\\alpha$, is an {\\it orientational} variable, which, being coupled to the internal {\\it shape} degrees of freedom, cannot be a constant of motion. In conclusion, in this non-Abelian case one has only two (instead of three  as in the Abelian case) commuting constants of motion, namely $S$ and $S^3$ (like in quantum mechanics).  This is also the outcome of the momentum map canonical reduction \\cite{mars,lib,con2} by means of adapted coordinates. Let us stress that $\\alpha$, $S^3$, $\\beta$ are a local coordinatization of any coadjoint orbit of SO(3) contained in the N-body phase space. Each coadjoint orbit is a 3-dimensional embedded submanifold and is endowed with a Poisson structure whose neutral element is $\\alpha$. This latter is also the essential coordinate for the definition of the {\\it flag} of spinors (see Refs. \\cite{payne} and \\cite{pauri1}, Section V). On the other hand, the spinor flag is nothing else than a unit vector $\\hat R$ orthogonal to $\\vec S$ \\cite{pauu}, which is going to be a fundamental tool in what follows.  By fixing non-zero values of the variables $S^3$, $\\beta = tg^{-1}{{S^2}\\over {S^1}}$ through second class constraints, one can define a (6N-8)-dimensional reduced phase space. However, the canonical reduction cannot be furthered by eliminating $S$, since $\\alpha$ is not a constant of motion.   C) The group-theoretical treatment of rigid bodies \\cite{lib} [Chapter IV, Section 10] is based on the existence of the realization of the (free and transitive) {\\it left} and {\\it right}  Hamiltonian actions of the SO(3) rotation group on either the tangent or cotangent bundle over their configuration space. Given a {\\it laboratory or space frame} ${\\hat f}_r$, the generators of the {\\it left} Hamiltonian action\\footnote{We follow the convention of Ref.\\cite{little}; note that this action is usually denoted as a {\\it right} action in mathematical texts. The $S^r_q$'s are the Hamiltonians, associated with the momentum map from the symplectic manifold to $so(3)^{*}$ [the dual of the Lie algebra $so(3)$],  which allow the implementation of the symplectic action through Hamiltonian vector fields.} are the non-Abelian constants of motion $S^1$, $S^2$, $S^3$, [$\\{ S^r,S^s \\} = \\epsilon^{rsu} S^u$], viz. the spin components in the {\\it space frame}. In the approach of Ref.\\cite{little} the SO(3) principal bundle is built starting from the {\\it relative configuration space} and, upon the choice of a body-frame convention, a gauge-dependent SO(3) {\\it right} action is introduced.  Similarly, taking into account the {\\it relative phase space} of any isolated system, one may investigate whether one or more SO(3) {\\it right} Hamiltonian actions could be implemented besides the global canonical realization of the SO(3) {\\it left} Hamiltonian action, which is a symmetry action. In other words, one may look for solutions ${\\check S}^r$, r=1,2,3, [with $\\sum_r ({\\check S}^r)^2 = \\sum_r (S^r)^2 = S^2$], of the partial differential equations $\\{ S^r, {\\check S}^s \\} =0$, $\\{ {\\check S}^r, {\\check S}^s \\} =- \\epsilon^{rsu} {\\check S}^u$ and then build corresponding {\\it left} invariant Hamiltonian vector fields. Alternatively, one may look for the existence of a pair ${\\check S}^3$, $\\gamma = tg^{-1} {{{\\check S}^2}\\over {{\\check S}^1}}$, of canonical variables satisfying $\\{ \\gamma ,{\\check S}^3 \\} =-1$, $\\{ \\gamma ,S^r \\} = \\{ {\\check S}^3, S^r \\} =0$ and also $\\{ \\gamma ,\\alpha \\} = \\{ {\\check S}^3, \\alpha \\} =0$. Local theorems given in Refs.\\cite{pauri2,pauri1} guarantee that this is always possible provided $N \\geq 3$. See Chapter IV of Ref.\\cite{lib} for what is known in general about the actions of Lie groups on symplectic manifolds. Clearly, the functions ${\\check S}^r$ do not generate symmetry actions, because they are not constants of the motion.  The inputs coming from A), B), C) together with the technique of the spin bases of  Ref.\\cite{lucenti} suggest the following strategy for the geometrical and group-theoretical identification of a privileged class of canonical Darboux bases for the N-body problem:  1) Every such basis must be a {\\it scheme B} (i.e. a canonical completion of {\\it scheme A})\\cite{pauri2,pauri1} for the canonical realization of the rotation group SO(3), viz. it must contain its invariant $S$ and the canonical pair $S^3$, $\\beta = tg^{-1}{{S^2}\\over {S^1}}$. Therefore, all the remaining variables in the canonical basis except $\\alpha$ are SO(3) scalars.  2) As said above, the existence of the angle $\\alpha$  satisfying $\\{ \\alpha , S \\} =1$ and $\\{ \\alpha ,S^3 \\} = \\{ \\alpha ,\\beta \\} =0$ leads to the geometrical identification of a unit vector $\\hat R$ orthogonal to $\\vec S$ and, therefore, of a orthonormal frame $\\hat S$, $\\hat R$, $\\hat S\\times \\hat R$, which will be called {\\it spin frame}\\footnote{The notation $\\hat {}$ means unit vector.}.  3) The study of the equations ${\\hat R}^2=1$ and $\\{ \\vec S\\cdot {\\hat R}, {\\hat R}^i \\} =0$ entails the symplectic result $\\{ {\\hat R}^i, {\\hat R}^j \\} =0$. As a byproduct we get a canonical realization of an E(3) group with generators $\\vec S$, $\\hat R{}{}$ [$\\{ {\\hat R}^i,{\\hat R}^j \\} =0$, $\\{ {\\hat R}^i, S^j \\} = \\epsilon^{ijk} {\\hat R}^k$] and fixed values of its invariants ${\\hat R}^2=1$, $\\vec S\\cdot \\hat R=0$  ( non-irreducible type 3 realization according to Ref.\\cite{lucenti}).   4) In order to implement a SO(3) {\\it Hamiltonian right action} in analogy with  the rigid body theory\\cite{lib}, we must construct an orthonormal triad or {\\it body frame} $\\hat N$, $\\hat \\chi$, $\\hat N\\times \\hat \\chi$. The decomposition  \\beq \\vec S = {\\check S}^1\\hat \\chi +{\\check S}^2 \\hat N\\times \\hat \\chi + {\\check S}^3 \\hat N \\,\\,\\, {\\buildrel {def}\\over =}\\,\\,\\, {\\check S}^r {\\hat e}_r, \\label{I0} \\eeq  \\noindent identifies the SO(3) scalar generators ${\\check S}^r$ of the right action provided they satisfy $\\{ {\\check S}^r,{\\check S}^s \\} = -\\epsilon^{rsu} {\\check S}^u$. This latter condition together with the obvious requirement that $\\hat N$, $\\hat \\chi$, $\\hat N\\times \\hat \\chi$ be SO(3) vectors [$\\{ {\\hat N}^r,S^s \\} =\\epsilon^{rsu}{\\hat N}^u$, $\\{ {\\hat \\chi}^r,S^s \\} =\\epsilon^{rsu}{\\hat \\chi}^u$, $\\{ {\\hat N\\times \\hat \\chi}^r,S^s \\} =\\epsilon^{rsu}{\\hat N\\times \\hat \\chi}^u$] entails the equations \\footnote{With ${\\check S}^r=\\vec S\\cdot {\\hat e}_r$, the conditions $\\{ {\\check S}^r,{\\check S}^s \\} = -\\epsilon^{rsu} {\\check S}^u$ imply the equations $\\vec S\\cdot {\\hat e}_r\\times {\\hat e}_s +S^iS^j \\{ {\\hat e}^i_r,{\\hat e}^j_s\\} = \\epsilon_{rsu} S^k{\\hat e}_k^u$, hence the quoted result.}   \\beq \\{ {\\hat N}^r,{\\hat N}^s \\} = \\{ {\\hat N}^r,{\\hat \\chi}^s \\} = \\{ {\\hat \\chi}^r,{\\hat \\chi}^s \\} =0. \\label{I1} \\eeq   To each solution of these equations is associated a couple of canonical realizations of the E(3) group (type 2, non-irreducible): one with generators $\\vec S$, $\\vec N$ and non-fixed invariants ${\\check S}^3=\\vec S \\cdot \\hat N$ and $|{\\vec N}|$; another with generators $\\vec S$, $\\vec \\chi$ and non-fixed invariants ${\\check S}^1=\\vec S\\cdot \\hat \\chi$ and $|{\\vec \\chi}|$. These latter contain the relevant information for constructing the angle $\\alpha$ and the new canonical pair ${\\check S}^3$, $\\gamma =tg^{-1}{{{\\check S}^2}\\over {{\\check S}^1}}$ of SO(3) scalars. Since $\\{  \\alpha , {\\check S}^3 \\} = \\{ \\alpha , \\gamma \\} =0$ must hold, it follows that the vector $\\vec N$ necessarily belongs to the $\\vec S$-$\\hat  R$ plane. The three canonical pairs $S$, $\\alpha$, $S^3$, $\\beta$, ${\\check S}^3$, $\\gamma$ will describe the {\\it orientational} variables of our Darboux basis, while $|\\vec N|$ and $|\\vec \\chi |$ will belong to the {\\it shape} variables. Alternatively, an anholonomic basis can be constructed by replacing the above six variables by ${\\check S}^r$ and three uniquely determined Euler angles $\\tilde \\alpha$, $\\tilde \\beta$, $\\tilde \\gamma$.   Let us consider  the case N=3 as a first example. It turns out that a solution of Eqs.(\\ref{I1}) corresponding to a {\\it body frame} determined by the 3-body system configuration only, as in the {\\it rigid body} case, is completely individuated  once two orthonormal vectors $\\vec N$ and $\\vec \\chi$, functions of the relative coordinates and independent of the momenta,  are found such that $\\vec N$ lies in the $\\vec S$ - $\\hat R$ plane\\footnote{Let us remark that any pair of orthonormal vectors $\\vec N$, $\\vec \\chi$ function only of the relative coordinates can be used to build a body frame. This freedom is connected to the possibility of redefining a body frame by using a configuration-dependent arbitrary rotation, which leaves $\\vec N$ in the $\\vec S$-$\\hat R$ plane.}. We do not known whether in the case N=3 other solutions of Eqs.(\\ref{I1}) exist leading to momentum dependent body frames. Anyway, our constructive method necessarily leads to momentum-dependent solutions of for Eqs.(\\ref{I1}) for $N\\geq 4$ and therefore to momentum-dependent or {\\it dynamical body frames}.   We  can conclude that in the N-body problem there are {\\it hidden structures allowing the identification of special dynamical body frames which, being independent of  gauge conditions, are endowed with a physical meaning}.  The following particular results can be proven:  i) For $N=2$, a single E(3) group can be defined: it allows the construction of an orthonormal {\\it spin frame} $\\hat S$, $\\hat R$, $\\hat R\\times \\hat S$ in terms of the measurable relative coordinates and momenta of the particles.    ii)  For $N=3$, $\\vec S = {\\vec S}_1+{\\vec S}_2$, a  pair of E(3) groups emerge associated with ${\\vec S}_1$ and ${\\vec S}_2$, respectively. In this case, besides the orthonormal {\\it spin frame}, an orthonormal {\\it dynamical body frame} $\\hat N$, $\\hat \\chi$, $\\hat N\\times \\hat \\chi$ can be defined such that ${\\check S}^1=\\vec S\\cdot \\hat \\chi$, ${\\check S}^2=\\vec S\\cdot \\hat N\\times \\hat \\chi$, ${\\check S}^3 = \\vec S\\cdot \\hat N$ are the canonical generators of a SO(3) Hamiltonian {\\it right} action. The non-conservation of ${\\check S}^r$ entails that {\\it the dynamical body frame evolves in a way dictated by the equations of motion}, just as it happens in the rigid body case.  It will be shown that for N=3 this definition of {\\it dynamical body frame} can be reinterpreted as a special global cross section ({\\it xxzz gauge}, where $x$ stays for $\\hat \\chi$ and $z$ for $\\hat N$; this outcome is independent from the particular choice made for $\\vec N$ and $\\vec \\chi$) of the trivial SO(3) principal bundle of Ref.\\cite{little}, namely a privileged choice of body frame. Actually, the three canonical pairs of orientational variables $S^3$, $\\beta = tg^{-1} {{S^2}\\over {S^1}}$; $S$, $\\alpha$; ${\\check S}^3$, $\\gamma = tg^{-1} {{{\\check S}^2}\\over {{\\check S}^1}}$, can be replaced by the anholonomic basis of three Euler angles $\\tilde \\alpha$, $\\tilde \\beta$, $\\tilde \\gamma$ and by ${\\check S}^1$, ${\\check S}^2$, ${\\check S}^3$ as it is done in Ref.\\cite{little}. In our construction, however, the Euler angles $\\tilde \\alpha$, $\\tilde \\beta$, $\\tilde \\gamma$  are determined as the {\\it unique set} of {\\it dynamical orientation variables}. Then, the remaining canonical pairs $q^{\\mu}$, $p_{\\mu}$, $\\mu =1,2,3$, of this spin-adapted Darboux basis describe the {\\it dynamical shape}  phase space.   While the above {\\it dynamical body frame} can be identified with the global cross section corresponding to the {\\it xxzz gauge}, all other global cross sections cannot be interpreted as {\\it dynamical body frames} (or {\\it dynamical right} actions), because the SO(3) principal bundle of Ref.\\cite{little} is built starting from the relative configuration space and, therefore, it is a {\\it static}, velocity-independent, construction. As a matter of fact, after the choice of the shape configuration variables $q^{\\mu}$ and of a space frame in which the relative variables have components $\\rho^r_a$, the approach of Ref.\\cite{little} begins with the definition of the body-frame components ${\\check \\rho}^r_a(q^{\\mu})$ of the relative coordinates, in the form $\\rho^r_a=R^{rs}(\\theta^{\\alpha}) {\\check \\rho}^s_a(q^{\\mu})$  \\footnote{$R$ is a rotation matrix, $\\theta^{\\alpha}$ are arbitrary gauge orientational parameters   and ${\\check \\rho}^r_a(q^{\\mu})$ is assumed to depend on the shape variables only and not on their conjugate momenta.}, and then extends it in a {\\it velocity-independent} way to the relative velocities ${\\dot \\rho}^r_a\\, {\\buildrel {def} \\over =}\\, R^{rs}(\\theta^{\\alpha}) {\\check v}^s_a$. In our construction we get instead $\\rho^r_a=R^{rs}(\\tilde \\alpha ,\\tilde \\beta ,\\tilde \\gamma ) {\\check \\rho}^s_a(q^{\\mu})$ in the {\\it xxzz gauge}, so that in the present case (N=3)  all {\\it dynamical} variables of our construction coincide with the {\\it static} variables in the xxzz gauge. On  the other hand, in  the relative phase space, the construction of the {\\it evolving dynamical body frame} is based on non-point canonical transformations      iii) For N=4, where $\\vec S = {\\vec S}_1+{\\vec S}_2+{\\vec S}_3$, it is possible to construct {\\it three} sets of {\\it spin frames} and {\\it dynamical body frames}  corresponding to the hierarchy of clusterings $((ab)c)$ [i.e. $((12)3)$, $((23)1)$, $((31)2)$] of the relative spins ${\\vec S}_a$  \\footnote{In a way analoguous to the angular momentum composition in quantum mechanics; note that this spin clustering is independent of the center-of-mass clustering associated with the Jacobi coordinates: the existence of these two unrelated clusterings might prove to be a useful and flexible tool in molecular physics.}. The associated three canonical Darboux bases share the three variables $S^3$, $\\beta$, $S$ (viz. $\\vec S$), while both the remaining three orientational variables and the shape variables depend on the spin clustering. This entails the existence of three different SO(3) {\\it right} actions with non-conserved canonical generators ${\\check S}^r_{(A)}$, A=1,2,3. Therefore, one can define three anholonomic bases ${\\tilde \\alpha}_{(A)}$, ${\\tilde \\beta}_{(A)}$, ${\\tilde \\gamma}_{(A)}$, ${\\check S}^r_{(A)}$ and associated shape variables $q^{\\mu}_{(A)}$, $p_{(A)\\mu}$, $\\mu =1,..,6$, connected by canonical transformations leaving $S^r$ fixed. These anholonomic bases and the associated {\\it evolving dynamical body frames}, however, have no relations with the N=4 {\\it static} non-trivial SO(3) principal bundle of Ref.\\cite{little}, which admits only local cross sections. As a matter of fact, one gets $\\rho^r_a=R^{rs}({\\tilde \\alpha}_{(A)} ,{\\tilde \\beta}_{(A)} ,{\\tilde \\gamma}_{(A)} ) {\\check \\rho}_{(A)a}(q^{\\mu}_{(A)}, p_{(A)\\mu}, {\\check S}^r_{(A)})$ instead of  $\\rho^r_a=R^{rs}(\\tilde \\alpha ,\\tilde \\beta ,\\tilde \\gamma ) {\\check \\rho}^s_a(q^{\\mu})$.  These results imply that, for N=4, the 18-dimensional relative phase space  admits three operationally well defined {\\it dynamical body frames}, and associated {\\it right} actions, and its coordinates are naturally splitted in three different ways into 6 dynamical rotational variables and 12 generalized dynamical shape variables. As a consequence, we get three possible definitions of {\\it dynamical vibrations}. Each set of 12 generalized dynamical canonical shape variables is obviously defined modulo canonical transformations so that it should even be possible to find local canonical bases corresponding to the local cross sections of the N=4 {\\it static} non-trivial SO(3) principal bundle of Ref.\\cite{little}.   Our results can be extended to arbitrary N, with $\\vec S = \\sum_{a=1}^{N-1}$ ${\\vec S}_a$. There are as many independent ways (say $K$) of spin clustering as in quantum mechanics. For instance for N=5, $K=15$ : 12 spin clusterings correspond to the pattern $(((ab)c)d)$ and 3 to the pattern $((ab)(cd))$ [$a, b, c, d = 1,..,4$]. For N=6, $K=105$: 60 spin clusterings correspond to the pattern $((((ab)c)d)e)$, 15 to the pattern $(((ab)(cd)e)$ and 30 to the pattern $(((ab)c)(de))$ [$a, b, c, d, e = 1,..,5$]. Each spin clustering  is associated to: i) a related {\\it spin frame}; ii) a related {\\it dynamical body frame}; iii) a related Darboux spin canonical basis with orientational variables $S^3$, $\\beta$, $S$, $\\alpha_{(A)}$, ${\\check S}^3_{(A)}$, $\\gamma_{(A)}=tg^{-1}\\, {{ {\\check S}^2_{(A)}}\\over {{\\check S}_{(A)}^1}}$, $A=1,.., K$ [their anholonomic counterparts are ${\\tilde \\alpha}_{(A)}$, ${\\tilde \\beta}_{(A)}$, ${\\tilde \\gamma}_{(A)}$, ${\\check S}^r_{(A)}$ with uniquely determined orientation angles] and  shape variables $q^{\\mu}_{(A)}$, $p_{\\mu (A)}$, $\\mu =1,.., 3N-6$. Furthermore, for $N \\geq 4$ we find the following relation between spin and  angular velocity: ${\\check S}^r = {\\cal I}^{rs}(q^{\\mu}_{(A)})\\, {\\check \\omega}^s_{(A)} + f^{\\mu}(q^{\\nu}_{(A)}) p_{(A)\\mu}$.  Let us conclude this Introduction with some remarks.  The $\\vec S=0$, {\\it C-horizontal}, cross section of the {\\it static} SO(3) principal bundle corresponds to N-body configurations that cannot be included in the previous Hamiltonian construction based on the canonical realizations of SO(3): these configurations (which include the singular ones) have to be analyzed independently since they are related to the exceptional orbit of SO(3), whose little group is the whole group.  While physical observables have to be obviously  independent  of the gauge-dependent {\\it static body frames}, they do depend on the {\\it dynamical body frame}, whose axes are operationally defined in terms of the relative coordinates and momenta of the particles. In particular, a {\\it dynamical} definition of {\\it vibration}, which replaces the $\\vec S=0$ {\\it C-horizontal} cross section of the {\\it static} approach\\footnote{Being connected to the Riemannian metric of the non-relativistic Lagrangian, this concept does not survive the transition to special relativity anyway.}, is based on the requirement that the components of the {\\it angular velocity} vanish. Actually, the angular velocities with respect to the dynamical body frames become now {\\it measurable quantities}, in agreement with the phenomenology of extended deformable bodies (e.g. the treatment of spinning stars in astrophysics).             In Section II the rest-frame description of N free particles together with the Jacobi normal relative coordinates is introduced and some further informations are summarized about the orientation-shape SO(3) principal bundle approach, both from the Lagrangian and the Hamiltonian point of view.  In Section III the {\\it canonical spin bases}, the {\\it spin frame} and the {\\it dynamical body frames} are introduced and the cases $N=2$, $N=3$ and $N\\geq 4$ are analyzed separately.  In Section IV a short account is given of N particles interacting through a potential.  Some final remarks are given in the Conclusions.  Appendix A contains the Lagrangian and Hamiltonian equations of motion in the {\\it static orientation-shape bundle} approach.  In Appendix  B the Lagrangian and Hamiltonian results of Subsection D of Section II are reformulated in arbitrary (but not Jacobi normal) relative coordinates.  In Appendix C detailed calculations are given for the case N=3.  In Appendix D some notions on Euler angles are reviewed.  In Appendix  E the gauge potential in the {\\it xxzz gauge} is evaluated.  In Appendix  F the construction of the {\\it canonical spin bases} is given for the $N=4$ case.    \\vfill\\eject ",
        "conclusions": "While it is possible to separate out absolute translations and relative motions in a global way, a global separation of absolute rotations from the associated notion of vibrations is not feasible because  of the non-Abelian nature of rotations.    This fact gave origin to the orientation-shape SO(3) principal bundle approach described in Ref.\\cite{little}. While the treatment of the non-relativistic ideal {\\it rigid body} is based on the notion of {\\it body frame} \\footnote{This schematization does not exist in special relativity: the only relativistic concept of rigidity are Born's rigid motions \\cite{born}.}, this and related concepts like angular velocities turn out  to be {\\it non-measurable gauge} quantities in this approach.  In this paper we have shown that, by using geometric and group-theoretical methods and by employing non-point canonical transformations, it is possible to introduce both a {\\it spin frame} and a N-dependent discrete number of {\\it evolving dynamical body frames} for the N-body problem, with associated {\\it dynamical orientation and shape variables}. Since all these concepts are introduced by means of a special class of canonical transformations defining the {\\it canonical spin bases}, they, and the related angular velocities, become {\\it measurable} quantities defined in terms of the particle relative coordinates and momenta. A notion of {\\it dynamical vibrations} is associated  to each spin basis. {\\it The spin bases are defined by a method based on the coupling of spins}, which does not require the use of Jacobi normal relative coordinates. However, the Hamiltonian is not the sum of the dynamical {\\it rotational} and {\\it vibrational} energies: this shows once again that rotations and vibrations cannot be separated in a unique global way.    Even if, in differential geometry, Darboux canonical bases have no intrinsic meaning, the physical requirement of the existence of the extended Galilei group with the spin invariant $S$ selects a preferred class of canonical coordinates adapted to the rotation subgroup SO(3). Moreover, various E(3) realizations emerge whose invariants have a physical meaning. The adaptation to these groups selects the final class of physically preferred Darboux canonical bases.  Our definition of {\\it dynamical body frames} is based on the existence of {\\it Hamiltonian non-symmetry right actions} of SO(3). We begin assuming that, like for rigid bodies, the body frame axes in the case N=3 be functions only of the particle relative coordinates.    For N=3, the orientation-shape SO(3) principal bundle is trivial and we recover its results in a {\\it xxzz gauge}. Interestingly we have found, however, that already for N=4 the {it dynamical body frame} axes depend also on the momenta so that any connection with the orientation-shape bundle approach is lost.  It is an open problem whether the use of more general body frames for N=3 and $N \\geq 4$ obtained by using the freedom of making arbitrary configurations dependent rotations (see footnote 9) may be used to simplify the free Hamiltonian and/or some type of interaction.    It is hoped that our results may be instrumental for nuclear, atomic and molecular physics, since a description based on spin clustering rather than on the standard Jacobi center-of-mass clusterings was lacking till now.  Let us observe that the extension of the {\\it dynamical body frames} to continuous deformable bodies (see Ref.\\cite{mate} for an initial study of the relativistic configurations of a Klein-Gordon field from this point of view) is a lacking piece of kinematical information and will be studied elsewhere.  The fact that we use non-point canonical transformations will make the quantization more difficult than in the orientation-shape bundle approach, where a separation of rotations from vibrations in the Schroedinger equation is reviewed in Ref.\\cite{little}. The quantizations of the original canonical relative variables and of the canonical spin bases will give equivalent quantum theories only if the non-point canonical transformations are unitarily implementable. These problems are completely unexplored.  In a future paper\\cite{iten1} we shall study the relativistic N-body problem, where the definition and separation of the center-of-mass motion are known to be a complicated issue.  This problem has found a solution within the Wigner-covariant rest-frame instant form of dynamics \\cite{lus}. It will be shown that concepts like reduced masses, Jacobi normal relative coordinates and tensor of inertia do not exist at the relativistic level. Yet, in the framework of the rest-frame instant form, both the oriantation-shape SO(3) principal bundle approach and the canonical spin bases can be defined just as in the non-relativistic case.  ACKNOWLEDGMENTS: We are warmly indebted to our friend Michele Vallisneri for his irreplaceable help in drawing our {\\it frames}.    \\vfill\\eject  \\appendix"
    },
    "0011/astro-ph0011423_arXiv.txt": {
        "abstract": "We present radial velocities for 144 globular clusters (GCs) around the Virgo gE NGC 4472 (M49), and ages and metallicities from co-added spectra of 131 GCs.  We confirm our earlier finding that the metal-poor GCs have a significantly higher velocity dispersion than the metal-rich GCs, and we find little or no rotation in the metal-rich GCs.  The velocity dispersion profile is consistent with isotropic orbits for the GCs and the NGC 4472 mass distribution inferred from X-ray data.  Our sample of GCs spans a metallicity range of $-1.6 \\leq [Fe/H] \\leq 0$ dex. The metal-poor and metal-rich GC populations are coeval within the errors, and all GCs older than 6 Gyr at 95\\% confidence. Our findings are consistent with a merger origin for NGC 4472, but other elliptical formation models cannot be ruled out. ",
        "introduction": "GC spectra were obtained in two observing runs, one in 1994 with WHT/LDSS-2 (Sharples et al. 1998), and the second in 1998 with CFHT/MOS (Beasley et al. 2000; Zepf et al. 2000). GC candidates for both runs were chosen from Washington photometry (Geisler et al. 1996), with colour ($0.5 < C-T_1 < 2.2$) and magnitude ($19.5 < V < 22.5$) selection. Spectra were obtained over the wavelength range 3600$-$6000 \\AA, with a resolution of 3$-$6 \\AA, and a velocity precision of 50$-$100 km/sec. Exposure times were 3$-$3.5 hours per mask. We have a total sample of 144 confirmed GCs in NGC 4472, out to $\\sim$ 7\\arcmin~ radius ($\\sim$ 30 kpc, 6 R$_{eff}$).  See Sharples et al., and Beasley et al. for more details. ",
        "conclusions": ""
    },
    "0011/astro-ph0011109_arXiv.txt": {
        "abstract": "We calculate analytically and numerically the Dyer-Roeder distance in perfect fluid quintessence models and give an accurate fit to the numerical solutions for all the values of \\ the density parameter and the quintessence equation of state. Then we apply our solutions to the estimation of $H_{0}$ from multiple image time delays and find that the inclusion of quintessence modifies sensibly the likelihood distribution of $H_{0}$ , generally reducing the best estimate with respect to a pure cosmological constant. Marginalizing over the other parameters ($\\Omega _{m}$ and the quintessence equation of state) , we obtain $H_{0}=71\\pm 6$ km/sec/Mpc for an empty beam and $H_{0}=64\\pm 4$ km/sec/Mpc for a filled beam. We also discuss \\ the future prospects for distinguishing quintessence from a cosmological constant with time delays. ",
        "introduction": "Quintessence (Caldwell et al. 1998) or dark energy is a new component of the cosmic medium that has been introduced in order to explain the dimming of distant SNIa \\ (Riess et al. 1998; Perlmutter et al. 1999) through an accelerated expansion while at the same time saving the inflationary prediction of a flat universe. The recent measures of the CMB at high resolution (Lange et al. 2000, de Bernardis et al. 2000, Balbi et al. 2000) have added to the motivations for a conspicuous fraction of unclustered dark energy with negative pressure. In its simplest formulation (see e.g. Silveira \\& Waga 1997), the quintessence component can be modeled as a perfect fluid with equation of state \\begin{equation} p_{q}=\\left( \\frac{m}{3}-1\\right) \\rho _{q}  \\label{eq stato q} \\end{equation} with $m$ in the range $0\\leq m<3$ ($m<2$ for acceleration). When $m=0$ we have pure cosmological constant, while for $m=3$ we reduce to the ordinary pressureless matter. The case $m=2$ mimicks a universe filled with cosmic strings (see e.g. Vilenkin 1984). More realistic models possess an effective equation of state that changes with time, and can be modeled by scalar fields (Ratra \\& Peebles 1988, Wetterich 1995, Frieman et al. 1995, Ferreira and Joyce 1998), possibly with coupling to gravity or matter (Baccigalupi, Perrotta \\& Matarrese 2000, Amendola 2000).  The introduction of the new component modifies the universe expansion and introduces at least a new parameter, $m$, in cosmology. Most deep cosmological tests, from large scale structure to CMB, from lensing to deep counting, are affected in some way by the presence of the new field. Here we study how a perfect fluid quintessence affects the Dyer-Roeder (DR) distance, a necessary tool for all lensing studies (Dyer \\& Roeder 1972, 1974). The assumption of constant equation of state is at least partially justified by the relatively narrow range of redshift we are considering, $% z\\in (0.4,2.5)$. We rederive the DR equation in quintessence cosmology, we solve it analytically whenever possible, and give a very accurate analytical fit to its numerical solution. Finally, we apply the DR solutions to a likelihood determination of $H_{0}$ through the observations of time-delays in multiple images. The dataset we use is composed of only six time-delays, and does not allow to test directly for quintessence; however, we will show that inclusion of such cosmologies may have an important impact on the determination of $H_{0}$ with this method. For instance, we find that $H_{0}$ is smaller that for a pure cosmological constant. In this work we confine ourselves to flat space and extremal values of the beam parameter $\\alpha $; in a paper in preparation we extend to curved spaces and general $\\alpha $. ",
        "conclusions": "\\bigskip  If 70\\% or so of the total matter content is filled by a new component with negative pressure and weak clustering, all the classical deep cosmological probes are affected in some way. Here we addressed the question of how the Dyer-Roeder distance changes when this new component, quintessence, is taken into account. We have shown that, particularly in the case of empty beam, the effect of the quintessence is to move the estimate of $H_{0}$ to lower values with respect to two standard models, and to increase the spread of the likelihood distribution. This is the first time, to our knowledge, that a full likelihood analysis of the almost entire set of time-delays available is performed. As a byproduct of our analysis, we produced fits for a large range of values of $\\Omega _{m}$ and $m$ accurate to within 5\\%$.$  The future prospects seem interesting for the time-delay method: with a not unrealistic increase in accuracy (or in the number of time delays), the quintessence could be detected and distinguished from a pure cosmological constant, thanks to the deepness to which lensing effects are observable.  An obvious improvement of our analysis, currently underway, is to investigate the dependence on $\\alpha \\neq 0,1$ and on curvature, producing a fully marginalized likelihood for $H_{0}$."
    },
    "0011/astro-ph0011279_arXiv.txt": {
        "abstract": "We numerically investigate the dynamics of how a close stellar fly-by encounter of a symmetrical circumstellar planetesimal disc can give rise to the many kinds of asymmetries and substructures attributed to the edge-on dusty disc of $\\beta$ Pic. In addition we present new optical coronagraphic observations of the outer parts of $\\beta$ Pic's disc, and report that the radial extent is significantly greater than was found in previous measurements. The northeasterly extension of the disc's midplane is now measured out to $1835$\\au~from the star; the southwesterly component is measured out to $1450$\\au. Thus the asymmetry in the length of the former with respect to the latter is approximately $25$ per cent. We proceed to use the length asymmetry induced in the distribution of simulation test particles as the principal diagnostic feature when modelling the disc response in order to constrain fly-by parameters. In particular we favour a low inclination prograde and near-parabolic orbit perturber of mass approximately $0.5$\\,M$_\\odot$. These initial conditions suggest that the perturber could have been physically associated with $\\beta$ Pic prior to the encounter. Thus we also consider the possibility that the perturber could be bound to $\\beta$ Pic: a consideration also of general interest where dust discs are known to exist in binary star systems. We show that a further consequence of a low velocity encounter is that the perturber could have captured planetesimals from the $\\beta$ Pic disc, and we deduce that as a result of this the perturber could display a dust disc that is presently amenable to observation. In some of our models, we can relate groupings of perturbed particles to the large-scale structure of the $\\beta$ Pic disc. The groupings correspond to: high eccentricity and inclination particles that reach apocentre and maximum height in the southwest, moderately eccentric and low inclination particles that reach apocentre in the northeast, and relatively unperturbed particles inside $\\sim$200\\au~radius. ",
        "introduction": "The starlight of $\\beta$ Pic scatters off dust grains in its edge-on circumstellar disc (Smith \\& Terrile 1984). Optical imaging reveals that the disc exhibits asymmetrical structure with respect an ideal disc on projected radial scales from $\\sim$50\\au~(Burrows et al. 1995, Mouillet et al. 1997) to $\\sim$1000\\au~(Smith \\& Terrile 1987, Kalas \\& Jewitt 1995). Kalas \\& Jewitt (1995) presented sensitive $R$-band images that showed the northeast (NE) extension could be detected out to $790$\\au, whereas the southwest (SW) extension was detectable out to just $650$\\au, corresponding to a $\\sim$20 per cent length asymmetry. More recently, the longer NE midplane extension has been discovered to carry several brightness enhancements from $\\sim$500 to $800$\\au~(Kalas et al. 2000; hereinafter KLSS), that have been attributed to a series of rings within the disc's midplane. Similar features in the SW extension were not detected and so eccentric rings were hypothesised.  KLSS presented a dynamical model in which the stellar fly-by scenario (Kalas \\& Jewitt 1995) for generating the disc asymmetries was applied to the $\\beta$ Pic system. Their finding was that the asymmetry types present in the $\\beta$ Pic disc have analogues in the dynamical response of a particle disc to a perturbing stellar fly-by encounter. Furthermore, they demonstrated that these features could occur simultaneously with the formation of eccentric circumstellar rings. A simple analysis of the perturbed disc model was found to be in good quantitative agreement with the observations of both $\\beta$ Pic's disc asymmetries and the midplane features; for particular values of the viewing angle and epoch, and fly-by parameters. In this paper we present new observations and measurements of the length asymmetry of $\\beta$ Pic's scattered-light disc, and then go on to simulate the dynamics of a stellar flyby encounter with an initially symmetrical disc of particles.  Similar numerical modelling to that presented here and in KLSS has been considered previously in the context of angular momentum transfer between the perturber's orbit and a protostellar disc (Clarke \\& Pringle 1993, Hall, Clarke \\& Pringle 1996), the dynamics of the Edgeworth-Kuiper belt (Ida, Larwood \\& Burkert 2000), and the tidal interaction of spiral galaxies (Toomre \\& Toomre 1972). Here we shall consider a wider range of encounter velocities (from bound orbits through to hyperbolic trajectories) and stellar masses (from M dwarfs to A stars, when scaled to the mass of $\\beta$ Pic) than has been studied before in any context. We also include a discussion of the orbital dynamics of captured particles with respect to the perturber. This paper can be considered as a companion to KLSS since here we also supplement their finding of circumstellar ring formation in similar simulations with an explanation of the dynamical origin of the features.  As in other works (Whitmire, Matese \\& Tomley 1988, Mouillet et al. 1997, KLSS) we assume that the simulation particles' distribution represents the distribution of an underlying disc of planetesimals, which are parent bodies whose infrequent and therefore dynamically inconsequential collisions supply the dust scattering surface that is observed in the real system. A more complete model should include consideration of grain removal and generation processes in determining the appearence of the perturbed disc (reviewed by Backman \\& Paresce 1993). However, we defer treatment of those issues to future work, and focus here on the first-stage problem of the dynamics. This stage is important in deducing the mass and orbital parameters of the postulated stellar perturber, for which it is possible to perform star catalogue searches (e.g. Kalas, Deltorn \\& Larwood 2000), and we proceed to do this by examining the length asymmetry, and other measures of the disc response, as a function of stellar encounter parameters. However, we choose to examine the length asymmetry as the principal diagnostic measure of the disc response since it is the most distinct such feature both in the observations and the dynamical models. In our future work we shall also consider the subtler asymmetries in greater detail.  In Section $2$ we present deep optical coronagraphic images of the disc. In Section $3$ we describe the results of our numerical investigations of the dynamics of tidally disrupted particle discs, and provide general discussion of the induced asymmetries and particle stripping. In Section $4$ we discuss our findings and summarise the conclusions. ",
        "conclusions": "Here we have considered the stellar fly-by hypothesis that may account for the asymmetrical structure in the scattered light dust disc of $\\beta$ Pic. We investigated the dynamics of the perturbed planetesimal disc in an attempt to learn more about the possible parameters of a fly-by encounter. Additionally, we have described the origin of transient circumstellar eccentric ring structures as being a general outcome of an encounter, due to the reflex motion of the primary star as the perturber passes through closest approach. The apparent ring system is actually an eccentric tightly-wound one-armed spiral pattern that gradually disappears owing to phase-mixing in the particles' orbital motion. In inclined-orbit simulations we find that the vertical disc response also has analogues in the scattered light disc of $\\beta$ Pic. We summarise the coplanar results in Fig. $18$, and the inclined-orbit results in Fig. $19$.  For the perturber we conclude that the mass should not be very different from $0.5$\\,M$_\\odot$ (corresponding to spectral type M0V). In this case we favour a low inclination prograde encounter with relative velocity of a few km\\,s$^{-1}$, which makes a bound perturber seem the most natural choice. One problem with this is that asymmetries would be destroyed in the second pericentre passage. Thus, given the deduced youth of the asymmetries, the perturber in this scenario is required to make a single close approach in the recent past (i.e. $\\sim$10$^5$\\yr~ago). This can be accounted for by making the perturber-$\\beta$ Pic system an initial wide binary that is perturbed by a massive passing field star $\\sim$10$^5$\\yr~in the past, such that the perturber is sent into a close approach orbit about $\\beta$ Pic. A more probable solution is to arrange for $\\beta$ Pic to have originally been a hierachical multiple system (cf. HD 141569; Weinberger et al. 2000) that became dynamically unstable on the age of the star, and with the result being that one component was sent into a disc-intercepting orbit. In any case the perturber should have captured planetesimals from the primary's disc and may exhibit a detectable scattered light disc as a result. A survey of a more extended field around $\\beta$ Pic than has been considered previously is in progress in order to test our hypotheses.  The dynamical response of the models suggests that the $\\beta$ Pic disc might be separable into different components corresponding to groupings of perturbed particle orbits. For example, we find that outwards of $\\sim$250\\au~the SW extension may not have a distinct midplane owing to pumping of particle inclinations. The positions close to pericentre of the moderately eccentric NE midplane particles could account for the measurement of a SW midplane out to $\\sim$650\\au, in superposition with the more diffuse vertically and radially extended component. When the optical image is vertically compressed (Fig. $2$) we are able to measure the SW extension to much larger radii ($1450$\\au). In our model the corresponding (very eccentric and inclined) particle orbits also intersect the disc near to their pericentres (down to $\\sim$50\\au) possibly accounting for the inferred inner warp of the disc, which is aligned with the outer flared envelope in the SW. In summary, there are three recognizable particle groupings that can be related to the morphology of the $\\beta$ Pic disc. These are: the highly eccentric and inclined particles that reach apocentre in the SW (extension B particles), the moderately eccentric and weakly inclined particles that reach apocentre in the NE (extension A particles), and the relatively unperturbed particles inside $\\sim$200\\au~radius (Figs. $12$ and $13$)."
    },
    "0011/astro-ph0011203_arXiv.txt": {
        "abstract": " ",
        "introduction": "The difference between the population of galaxies in the field and in clusters is well known since the early 1930s. Oemler (1974) analyzed the proportion of elliptical, S0 and spiral galaxies and defined three different types of clusters: spiral rich, spiral poor and cD clusters. Melnick and Sargent (1977) showed that the inner region of clusters are typically populated by ellipticals and spirals are predominating in the outer regions (the well known morphological segregation, MS).  In a classic work, Dressler 1980b analyzed a sample of 55 nearby clusters finding a correlation between the fraction of different morphological types T and the local projected galaxy density $\\Sigma$ (hereafter T-$\\Sigma$ relation). Dressler concludes that galaxy morphology is a function of the local clustering rather than global conditions related to the cluster environment.  Samrom$\\grave{a}$ and Salvador Sole (1990) have consider a test to explore the nature of the morphological segregation. Using Dressler's data, artificial clusters were generated by randomly repositioning the polar coordinates of galaxy members around the center of the clusters. In this way  any subclustering present is erased maintaining the radial clustercentric distance of each galaxy so that the global cluster properties are conserved. The T-$\\Sigma $ relation is analyzed  for the randomized galaxies in the clusters and compared to the results for the real galaxy positions. The results of these tests are indistinguishable indicating that the T-$\\Sigma $ relation is controlled by global conditions rather than local subclustering.  Whitmore et al. (1991,1993, hereafter WGJ ) re-examined Dressler's sample of galaxies in clusters and suggested that the morphology-clustercentric distance relation is more fundamental than the morphology-local galaxy density relation. These authors  use the clustercentric normalized distance as the independent parameter, and compare the morphological fractions at the same normalized clustercentric radii but with different values of the local galaxy density. WGJ use a \"characteristic radius\", ${r_{opt}}$, as the radius at which the cumulative projected galaxy density falls bellow 20 galaxies $Mpc^{ -2}$. These authors find that for small radii (about 0.5 Mpc) the elliptical fraction rise rapidly. On the other hand, the spiral fraction is essentially zero at the cluster center. WGJ interpret these results as an indication that cluster center conditions play a  key role in determining galaxy morphologies in clusters and suggest that a destructive rather than a formation mechanism may be controlling their relative fractions.  The question whether local galaxy density or radial distance from the cluster center is best correlated with morphological types is still open. Dressler et al. (1997) reanalyzed the morphology-density relation by dividing their sample into centrally concentrated, regular clusters; and low-concentration, irregular clusters, finding a similar T-$\\Sigma$ relation in both subsamples.  Nevertheless, they find a significant excess of SO's and a smaller spiral fraction in the centrally concentrated clusters when compared to the low-concentration systems.  Dressler et al. (1997, hereafter D97) also analyzed a sample of 10 clusters at redshifts between 0.37 and 0.56, and derive the corresponding T-$\\Sigma$ relation. For these  distant clusters the authors find that the fraction  of S0 galaxies is 2-3 times smaller than in low-redshift clusters suggesting that S0's are generated in large numbers only after cluster virialization.  It should be recalled, however, the possible presence of systematic effects related to projection biases in the selection of clusters from two dimensional data (see for instance Valotto, Moore, and Lambas, 2001), which may cause an artificial increase of the late type fraction in distant clusters.  Several mechanisms have been proposed in order to explain the morphological segregation. Among them we can mention ram pressure stripping (Gunn and Gott 1972, Abadi et al. 1999), gas evaporation (Cowie and Songaila 1977), merging (Lavery and Henry 1988), galaxy harassment (Moore et al. 1996), tidal striping (galaxy-galaxy and from mean cluster field; Bird and Valtonen 1990), tidal shaking (galaxy-galaxy and from the mean cluster field; Miller 1988), galactic cannibalism (Ostriker and Tremaine 1975), truncated star formation (Larson et al. 1980), etc. Also, it should be considered the effects provided by the initial conditions which may also play a role in the morphological segregation. In hierarchical models of structure formation such as CDM, different scales become non-linear simultaneously, so that the initial conditions of galaxy formation in cluster and group environments may differ from the field (intergalactic medium, tidal effects, mergers, etc).    In this work we study the relative fraction of galaxy morphologies in clusters  as a function of global cluster parameters (projected mass and gas density profile), and the local projected galaxy density. Our analysis is aimed to provide a better understanding of the relevance of the proposed mechanisms involved in morphological segregation (MS) in clusters given their different dependence on gas, mass and galaxy content.  The paper is organized as follows. In Section 2 we briefly review the data used for this study. In Section 3 we explore the MS as a function of different global and local parameters, and perform a simple test to compare their relative importance. Section 4 provides a discussion of the  main results and some implicances for  galaxy evolution in clusters. We include in Appendix A a short discussion on the corrections applied to deal with observational biases. ",
        "conclusions": "In hierarchical scenarios for structure formation in the universe, clusters are assembled from smaller subunits, so that the local galaxy density, associated to these primordial clumps, could play a significant role in the segregation of morphological types. On the other hand, the several mechanisms related to galaxy evolution within the potential well of clusters (RAM pressure, tidal effects, etc.) could also explain the morphological segregation and its relation to global parameters.   The results of our analysis suggests that there are different mechanisms controlling the morphological segregation depending  on  the galaxy  environment. We find that mechanisms of global nature dominate in high density environments, namely, the virialized region of clusters,  while local galaxy density as defined by Dressler 1980b is the relevant parameter in the outskirts where the influence of the cluster as a whole is relatively small compared to local effects.  As it can be inferred by inspection to figure 4, a primary parameter in the segregation of morphological types in high density regions is the global cluster mass density at the position of the galaxies, computed using the scaling relationship between the mean cluster temperature  and the projected  NFW mass density profile. We find that the relative fraction of galaxy morphologies shows a stronger dependence on the local NFW projected density profile compared to other radial distances normalizations such as $r_{500}$ or $r_{opt}$. Therefore, these results might be applied to other studies of galaxy properties in clusters such as star formation rate, fraction of blue galaxies, etc. that show a significant dependence on radial distance to the cluster centers.  By comparison of the results corresponding to the NFW and the $\\beta$ model profiles (projected mass and gas density respectively) we conclude that the relative fraction of galaxy morphologies is more strongly correlated with the NFW profile. This result may serve to assess the importance of tidal effects and gaseous phenomena operating in the transformation of spiral galaxies into S0s.  Our results give support to the idea that tidal force effects produced by the cluster potential, galaxy harassment, or truncated star formation among others physical mechanisms would be primary in driving the observed morphological segregation  in clusters of galaxies.  We have also  taken into account different cluster properties in our analysis by considering subsamples of high/low cluster temperature and high/low X-ray luminosity. The results of these analysis are  similar to those obtained for the total sample although in the high density regions of hot clusters we find a tendency of a stronger morphological segregation as a function of $\\Sigma_{Mass}$."
    },
    "0011/gr-qc0011072_arXiv.txt": {
        "abstract": "{} We have developped a procedure for the search of periodic signals in the data of gravitational wave detectors. We report here the analysis of one year of data from the resonant detector Explorer, searching for pulsars located in the Galactic Center (GC). No signals  with amplitude greater than $\\overline{h}= 2.9~10^{-24}$, in the range 921.32-921.38 Hz, were observed using data collected over a time period of 95.7 days, for a source located at $\\alpha=17.70 \\pm 0.01$ hours and $\\delta=-29.00 \\pm 0.05$ degrees. Our procedure can be extended for any assumed position in the sky and for a more general all-sky search, even with a frequency correction at the source due to the spin-down and Doppler effects. ",
        "introduction": "Periodic or almost periodic gravitational waves (g.w.) are emitted by various astrophysical sources. They carry important information on their sources (e.g., spinning neutron stars, accreting neutron stars in binary systems) and also on fundamental physics, since their nature can test the model of General Relativity \\cite{Schutz_sigrav,thorne}. The main feature of continuous signals which allows them to be detected is that, despite the weakness of the signal (compared to typical amplitudes for bursts), it is possible to implement procedures that build up the signal to noise ratio (SNR) in time. The natural strategy for searching for monochromatic waves is to look for the most significant peaks in the spectrum. In this case the SNR increases with the observation time $t_{obs}$. In fact, as $t_{obs}$ increases, the frequency resolution of the spectrum also increases - the frequency bin gets smaller $\\delta\\nu=\\dps{1\\over t_{obs}}$ - thus the content of the noise power in each bin decreases with $t_{obs}$, while the signal is not dependent on the length of observation time\\footnote{The noise power {\\it{density}}, i.e. the power spectrum, is independent of $t_{obs}$, while the signal energy {\\it{density}} grows with $t_{obs}$}. For a periodic signal of amplitude $\\bar{h}$ at the frequency $\\bar \\nu$ the squared modulus of the Fourier Transform provides ${\\bar{h}^2}$ with a noise contribution of $2~S_h(\\bar{\\nu})~\\delta\\nu$, being $S_h(\\bar{\\nu})$ the two-sided noise power spectrum of the detector (measured in Hz$^{-1}$).   Thus the SNR for periodic signals is: \\beq SNR={{\\overline{h}^2 t_{obs}} \\over{2 S_h(\\overline{\\nu})}} \\label{snr_tuttonoto} \\eeq  Eq.\\ref{snr_tuttonoto} holds if the instantaneous frequency of the continuous signal at the detector is  known. The analysis procedure in this case is ``coherent'', since the phase information contained in the data is used and the sensitivity (in amplitude) increases with the square-root of the time. However in some cases it may be impossible, for various reasons (see later in sect. 3.1), to perform a single Fourier Transform over all the data. This means that the observation time has to be divided in $M$ subperiods, such that the spectral resolution of the spectra becomes $\\delta \\nu ^{'}=M / t_{obs}$ and the corresponding $SNR$ is $M$ times smaller than that given by eq. \\ref{snr_tuttonoto}.  The $M$ spectra can be combined together by incoherent summation, that is by averaging the square modulus. In this case the final spectral resolution is again $\\delta\\nu ^{'}$ but there is still some gain as this procedure reduces the variance of the observed noise in the bin. We obtain: \\beq {SNR}^{\\prime}={{\\overline{h}^2 t_{obs}} \\over{2 S_h(\\overline{\\nu}) \\sqrt{M}}} \\label{snr_nontuttonoto} \\eeq   In general, if the signal is monochromatic but frequency modulated due to the detector-source relative motion, processing techniques exist which can recover the sinusoidal case if the source direction is known. One of the standard ways of detecting such signals is through appropriate  resampling of data, better known in the radio astronomy comunity (where this technique is commonly used) as ``data stretching'' (see for example \\cite{camilo1}). In the case of radio pulsar searches, the location of the source is usually known (the data come from a radio telescope pointing in a particular direction) but some parameters of the system need to be estimated and this is done by a ``timing solution which is phase coherent over the whole data set''\\cite{camilo1}.  However for gravitational waves, especially when searching a large parameter space, it is doubtful that strategies developed for radio pulsar searches can simply be adapted: the expected low SNR values for g.w. signals really modifies the nature of the search strategies that can be employed. In recent ms pulsar searches, for example in \\cite{camilo1}, the signal is strong enough to allow suspected pulsars to be identified by {\\it visual} inspection of the results of the final stages of the analysis procedure.    For gravitational waves the study of the implementation of optimum analysis procedures is still in progress \\cite{brady,map,krolak,aroma,brady2}.  In the present analysis, we analyzed the data of the Explorer detector to search for ms pulsars located in the GC, assuming their intrinsic frequency to be constant over the analysis time.  The procedure we used in this study relies on a data base of FFTs, computed from short stretches of data (short with reference to the effects of the Doppler shift, as will be described later in this paper). These short FFTs are then properly combined together to provide a new set of FFTs with higher frequency resolution, representing the signal in the frequency range selected for the study.   The paper is organized as follows: in section \\ref{expl} we briefly review the characteristics of the detector during the 1991 run, in section \\ref{main} we describe the main aspects of the procedure; in section \\ref{result} we present the results obtained. The Appendices clarify some aspects of the analysis procedure and discuss the extent to which the constraints that we have introduced to describe the procedure can be relaxed in order to account for different sources. ",
        "conclusions": "\\label{final} A first analysis of the data obtained with the Explorer detector in 1991 was performed, with the aim of searching for continuous g.w.. The analysis was limited to the frequency range 921.32-921.38 Hz, which contained the plus resonance of the detector, where sensitivity was highest. Doppler corrections on the g.w. frequency were made under the assumption that the source was still in the GC without any intrinsic frequency spin down.  No signals were observed with amplitude greater than $\\overline{h}= 2.9~10^{-24}$, using data collected over 95.7 days, for a source located at $\\alpha=17.70 \\pm 0.01$ hours and $\\delta=-29.00 \\pm 0.05$ degrees, having a spindown parameter $\\tau \\ge 10{^8}$ years (that is $\\dot{p} \\le 1.7 \\cdot 10^{-19}$ s/s).  The procedures adopted here can be applied to any assumed position in the sky of a g.w. source, for a greater frequency range, or even for a frequency correction at the source due to spin down and Doppler effects.    \\par \\newpage \\appendix"
    },
    "0011/astro-ph0011035_arXiv.txt": {
        "abstract": "{ We present the high resolution spectrum of the isolated neutron star \\rxj, obtained with the Reflection Grating Spectrometer on {\\it XMM-Newton}, complemented with the broad band spectrum observed with the EPIC PN camera. The spectrum appears smooth, with no evidence for strong photospheric absorption or emission features. We briefly discuss the implications of our failure to detect structure in the spectrum. ",
        "introduction": "There is now compelling evidence that the {\\it ROSAT} source RXJ0720.4-3125, first discovered in a systematic survey of the galactic plane (Haberl et al. 1997), is a nearby isolated neutron star. Its spectrum appears very soft ($kT_{\\rm e} \\sim 80$ eV), and there is little attenuation of the X-ray flux by interstellar absorption. There is a plausible optical counterpart (Motch \\& Haberl 1998; Kulkarni \\& van Kerkwijk 1998), and the measured X-ray to optical flux ratio basically rules out an accreting (low-mass) X-ray binary, or an isolated white dwarf. Perhaps the strongest piece of evidence is the detection of a fairly shallow, quasi-sinusoidal modulation of the X-ray flux, at a period of 8.39 sec (Haberl et al. 1997). The stability of this period, as measured with {\\it ROSAT} over the period 1993-1996, suggests that we are seeing modulation by the spin of a compact object, most likely a neutron star.  The sinusoidal shape of the modulation in turn suggests that it is due to changing visibility, and possibly limb-darkening effects, of a hot polar cap (or polar caps), against a cooler atmosphere. Such a region could be heated by accretion from the interstellar medium, with the accretion stream focused by a magnetic field of order $B_{\\rm surface} \\lax 10^{10}$ G (the field can be no stronger, or else it would prevent accretion onto the surface, given the spin period and likely values for the density of the insterstellar medium). Alternatively, the object could be releasing heat from decay of a very strong magnetic field, which would at the same time modulate the heat flow through the atmosphere due to the influence of the field on the thermal conduction in the direction transverse to the field (Greenstein \\& Hartke 1983).   These issues can be directly addressed by studying both the spin history, as well as the detailed photospheric emission spectrum of the object. With {\\it XMM-Newton}, we have recently obtained data of superior sensitivity that will allow us to pursue both these possibilities.  In this {\\it Letter} we present a preliminary analysis of the high-resolution spectrum of \\rxj, obtained with the Reflection Grating Spectrometer (RGS; den Herder et al. 2001), which covers the range 5-35 \\AA\\ at an approximately constant resolution $\\Delta\\lambda = 0.05$ \\AA. Data were obtained simultaneously with the EPIC focal plane imaging cameras (Turner et al. 2001), of which the PN camera (Str\\\"uder et al. 2001) offers the combination of high sensitivity and high time resolution. This latter dataset will be used to investigate the evolution of the spin period. Unfortunately, however, we currently lack the satellite dynamical data to perform an accurate barycentric correction to the photon arrival times, and we defer an accurate measurement of the spin period to a forthcoming paper. ",
        "conclusions": "We present the atmospheric spectrum of the isolated neutron star \\rxj, as measured with the RGS and EPIC-PN instruments on {\\it XMM-Newton}. We find that the spectral shape is smooth, without evidence for strong spectral features. The absence of electron or proton cyclotron resonances in the RGS band would appear to exclude average surface magnetic field strengths of $B\\simeq (0.3-2.0)\\times 10^{11}$ and $(0.5-2.0)\\times 10^{14}$~G. The implications of the apparent absence of atomic transitions are strongly model dependent, and we can therefore not rule out either a pure hydrogen atmosphere or a metal-rich atmosphere, if the magnetic field strength satisfies certain restrictions."
    },
    "0011/astro-ph0011345_arXiv.txt": {
        "abstract": " ",
        "introduction": "Warps of disk galaxies are a common phenomenon (as common as spirtal structure), yet they are still not fully understood. In this review I will try to summarize their observed properties, including preliminary results of a new HI survey we have carried out with Westerbork, and try to relate these to several proposed mechanisms for explaining warps. I will not attempt a historic review of the theoretical work on this subject; such reviews can be found in Binney (1992), or Kuijken (1998), for example. ",
        "conclusions": "In spite of being known now for many decades, warps are still a puzzle. The observational situation and theoretical models to explain the data are still evolving. In this talk I have tried to summarize the situation as follows.  \\begin{enumerate} \\item Warps are a very common phenomenon. However, in addition to the classic `grand-design' integral-sign warps, in many cases the warp is asymmetric or even one-sided. Models for warps should therefore not only address the symmetric, regular ones. Quite possibly there is a good analogy here to spiral structure, which also displays varying degrees of regularity. \\item Normal modes (equilibrium configurations of a tilted, warped disk precessing about the symmetry axis of a flattened halo) interact strongly with the dark halo. As a consequence they are easily and strongly damped, or, in special circumstances, excited. \\item Satellite tides are generally too weak to produce warps of the amplitudes observed. However the dark halo can under certain circumstances respond to the satellite in a way that significantly adds to the tidal perturbation on the disk. A useful diagnostic for satellite tides as explanation for a warp is the orientation of the warp with respect to the satellite orbit. In the case of the LMC/Galaxy system, this orientation appears to be different from predictions. \\item Accretion of material can generate warps through a continual change in the orientation of the halo symmetry plane, to which different parts of the disk respond on different timescales. Such models are yet to be investigated in full detail; in particular the question of whether the observed frequency of asymmetric warps can be reproduced by this kind of model may be a useful avenue to explore. \\end{enumerate}"
    },
    "0011/astro-ph0011459_arXiv.txt": {
        "abstract": "We report performance verification observations of the giant elliptical galaxy M87 in the Virgo Cluster with the MOS, pn, and optical monitor instruments on board of XMM-Newton. With the energy sensitive imaging instruments MOS and pn we obtain the first spatially constrained X-ray spectra of the nucleus and the jet of the galaxy. The good photon statistics of the pn and MOS allow a detailed analysis of the radial temperature and abundance distribution of 6 elements. The data provide no indication of a multi-temperature structure for radii $\\ge $2$'$. An apparent sharp metal abundance drop deduced for the regions inside this radius is probably due to resonant line scattering. ",
        "introduction": "The giant elliptical galaxy M87 as the central dominant galaxy of the northern part of the Virgo cluster (e.g. Binggeli et al. 1987) displays a large variety of very interesting astrophysical phenomena (see e.g. R\\\"oser \\& Meisenheimer 1999). Because M87 is so bright and close, 17 - 20 Mpc (Freedman et al. 1994, Tammann \\& Federspiel 1997), these astrophysical processes can be studied in M87 in more detail than in any other comparable object.   M87 (Virgo A, 3C 274) is one of the first known extragalactic radio sources and its halo is the first extragalactic X-ray source to have been identified (Bolton \\& Stanley 1948, Byram et al. 1966). The highly peaked surface brightness profile of the extended X-ray emission subsequently observed with EINSTEIN implied an X-ray gas cooling time considerably smaller than the Hubble time, leading to the suggestion that M87 harbours a significant ``cooling flow'' with a mass condensation rate of $\\sim 10$ M$_{\\odot}$ yr$^{-1}$ (Stewart et al. 1984, Fabian et al. 1984). EINSTEIN Crystal Spectrometer and Solid State Spectrometer observations of M87 provided X-ray line data which lent further support to the existence of a multi-temperature plasma in the core of the M87 halo, as expected for a cooling flow model (Canizares et al. 1979, 1982, Mushotzky \\& Szymkowiak 1988).  M87 also houses an active nucleus with a one-sided jet observed at radio (e.g. Owen 1989), optical (e.g. Sparks et al. 1996), and X-ray wavelengths (Schreier et al. 1982, Neumann et al. 1997, Harris et al. 1997, 1999). The power output of the jet and the unseen counterjet obviously feeds a complex system of inner ($r \\le 25$ kpc) and outer radio lobes ($r \\le 40$ kpc, e.g. Owen et al. 2000) which partly interact with the thermal gas giving rise to distinct features in the X-ray surface brightness distribution (e.g. Belsole et al. and references therein). The ability to obtain sensitive spatially resolved X-ray spectroscopy with the XMM-Newton observatory offers the opportunity to gain new insights into these phenomena.  This paper reports on the analysis of the XMM-Newton (Jansen et al. 2000) performance verification phase observation of M87 including results from the MOS (Turner et al. 2000), pn (Str\\\"uder et al. 2000) and Optical Monitor (Mason et al. 2000) instruments. An analysis of the morphology and spectroscopy of the X-ray surface brightness enhancements related to the radio lobes of M87 is presented in an accompanying paper by Belsole et al. (2000), and the results from high resolution spectroscopy with the RGS instrument are discussed in a forthcoming paper by Sakelliou et al. (2000). ",
        "conclusions": "In Fig. 7 the spectral energy distribution of the bright knots in the jet are compared with a model for synchrotron radiation (from an electron spectrum of slope 2.3, with $\\gamma_{\\rm min} = 100$ and $\\gamma_{\\rm max} = 2 \\times 10^6$ in the equipartition magnetic field of 265 $\\mu $Gauss). Also shown is the contribution of self-Comptonization for an equipartition field. Both components fall short of explaining the flux observed with XMM-Newton. The observations could be matched by synchrotron emission from an electron spectrum which extends to higher energies above a break at $\\gamma \\sim 8 \\times 10^5$, but the required steepening of the electron spectral slope of $\\sim 1.5$ is not naturally predicted by simple models. The high electron energies and consequently short lifetimes imply that the electrons are accelerated in-situ at the X-ray knot (see also Heinz \\& Begelman 1997). Lifetime effects could then explain why the optical and X-ray knots appear in slightly different positions.  The good agreement of the temperature profiles between the different XMM-Newton instruments as well as the ASCA and BeppoSAX instruments is a sign that the overall calibration of the detector response is already very good.  One of the surprising results is the non-detection of good evidence for a multi-temperature cooling flow structure for which M87 was expected to be a good target. We also do not find the signature of oxygen absorption claimed to be seen by Buote (2000) in ROSAT PSPC data.  A striking result is the sharp drop of all deduced metal abundances at radii smaller than about 1 arcmin. This effect can be observed directly in Fig. 8. We have shown that resonant line scattering is important and could be responsible for this effect. This effect has also to be taken into account in the interpretation of XMM-RGS spectra of nearby clusters, since some of the relevant emission lines may suffer from scattering out-off the field-of-view of the spectrometer.  The decrease of the abundance profiles in the outer regions show a significant trend with element mass. The decrease from $r = 1.5'$ to $r = 10'$ is about a factor of 2 or larger for Fe, S, Ar, less for Si (a factor of $\\sim 1.7$), and the profile is almost constant for O. This can be explained by a contribution to a homogeneous distribution by SN type II, the main source for O, and a SN type Ia yield of primarily heavy elements which is more concentrated in the center of the M87 halo."
    },
    "0011/astro-ph0011490_arXiv.txt": {
        "abstract": "We present a framework for the study of muon density patterns on the ground due to showers produced in the atmosphere by cosmic rays incident at high zenith angles. As a checking procedure predictions of a model based on such a framework are compared to simulation results. ",
        "introduction": "The old idea to detect neutrinos through Horizontal Air Showers (HAS) \\cite{Berezinskii} has been recently brought to attention with the calculation of the acceptance of the Auger Observatories for the detection of high energy neutrinos \\cite{Capelle_1}. At large zenith angles, cosmic rays (whether they are protons, heavier nuclei, or even photons) develop ordinary showers in the top layers of the atmosphere in a very similar fashion to the well understood vertical showers. Their electromagnetic component is however almost completely absorbed by the greatly enhanced atmospheric slant depth and prevented from reaching ground level. By contrast, high energy neutrinos may also induce Horizontal Air Showers (HAS) deeper into the atmosphere and hence resemble vertical air showers in their particle content at ground level, provided they interact sufficiently close to an air shower array.  The main background to neutrino induced HAS of very high energy (in the EeV range and above) is expected to be due to the remaining muon component of the cosmic ray showers, after practically all the electromagnetic component is absorbed. A preliminary study of the air shower data at zenith angles above $60^o$ from the Haverah Park array has shown that the observed rate is consistent to that expected from the ultra high energy ($E > 1$~EeV) cosmic rays arriving at large zenith angles \\cite{Ave_1}. These {\\sl muon showers} that penetrate the whole atmospheric slant depth to reach ground level are the object of this study which is intended to be detector independent. As the neutrino induced HAS rate that could be expected is quite low, zenith angle misreconstructions of vertical showers can also constitute a serious potential background to neutrino detection. However the relative importance of this background is very dependent on technical details of the experimental setup and, in particular, on the angular resolution actually achieved. Thus these matters cannot be further addressed in this article.  The interest in HAS induced by ordinary cosmic rays is not just limited to understanding the background of neutrino induced showers. The measurement of high zenith angle showers would enhance the aperture of existing air shower array data and it would increase the data on cosmic ray arrival directions to previously inaccessible directions in galactic coordinates \\cite{Ave_1}. Besides these obvious advantages, high zenith angle cosmic ray showers are unique because they select the muon component of the showers and they naturally probe a region of phase space which is quite inaccessible from vertical air shower measurements: the shower core.  The cosmic ray induced HAS are different mainly because they consist of muons which are produced far away from ground level. The particle density profiles for HAS induced by protons or heavy nuclei display complex muon patterns at the ground which result from the long path lengths travelled by the muons in the presence of the Earth's magnetic field \\cite{Hillas,Andrews}. These complex patterns are difficult to analyse \\cite{Dedenko} because the conventional approaches used for interpretation of low zenith angle showers ($< 60^\\circ$) are usually based on the approximate circular symmetry of the density profiles. The analysis of HAS produced by cosmic rays requires a radically different approach.  Here we present a qualitative explanation of the ground muon patterns as well as analytical expressions for the muon density profile based on simple ideas. Our approach agrees remarkably well with simulation results for a wide range of distances to the core and zenith angles between $60^\\circ$ and $90^\\circ$. Besides its convenience for data analysis which would otherwise depend on lengthy and cumbersome simulation results, the model gives a lot of insight into HAS and their complementarity to vertical showers and can be used for calculating the expected HAS rates for future detectors such as the Auger Observatories \\cite{Ave_prep}. We also briefly discuss a simple method involving rotations for approximately generating muon ground density profiles for any azimuthal incidence using a single shower of the same zenith angle.  The article is organized as follows. In section \\ref{toy} we present a toy model which describes the qualitative behavior of the ground muon patterns. In section \\ref{geomag} we discuss quite generally the implications of geomagnetic fields and the geometry implied. In section \\ref{geotoy} we extend the toy model to account for the magnetic field. In section \\ref{param} we give a detailed prescription to calculate muon densities for any azimuthal and zenith angle. In section \\ref{asymmetries} we discussed the asymmetries that arise at ground level and introduce a number of ways to account for them. In section \\ref{results} we compare the analytical densities obtained to those obtained with Monte Carlo simulations and finally in section \\ref{summary} we summarize the achievements, briefly enumerating possible applications. ",
        "conclusions": "\\label{summary}  We have successfully obtained a framework for studying the muon content of HAS produced by protons in the atmosphere. It can be effectively considered as a semianalytical parameterization for the muon number densities in the transverse plane for HAS. Although it relies on Monte Carlo it can be considered as an alternative to it since the simulation is only needed in first instance to obtain the inputs ignoring the magnetic effects and hence eliminating azimuthal angle dependences. Although only proton showers of $10^{19}~$eV energy have been considered for the Haverah Park site, this framework is completely general and can be applied to showers of any energy and composition located in any site in a completely analogous fashion.  The method consists of splitting the muon propagation in a shower into two fictitious parts, one which takes place in the absence of a magnetic field and a second one in which the magnetic effects are separately analysed. As the first part preserves the approximate circular symmetry of the shower development it allows simple parameterizations in terms of a single variable. The second part can be implemented in an analytical way. As a spinoff we have obtained a way to obtain ground density profiles of fixed zenith angle but any azimuthal angle from a single shower of the corresponding zenith.  We expect to refine this model by improving the fits for the required inputs and including ignored effects such as correlations between muon impact position and distance to production site, muon decay, and distributions of distance to production site. These ignored effects could be implemented in an analogous way to the introduction of the energy distribution. Although these refinements may be necessary in the future, we do not feel at this time there is need to be so accurate since there are many uncertainties in the simulation programs which surely exceed the precision level of the results obtained. The simulations used require extrapolations of cross sections and distributions of particle interactions in energy regions which are well above those explored in current particle physics experiments or those to be made in the near future.  Besides giving insight into the complexity of large zenith angle showers, the results presented here have a rather broad range of applications connected with the detection of Horizontal Air Showers. The need of systematic simulations of complicated high energy showers which is very time consuming can be avoided. Such a task is of great help for the analysis of large zenith angle data as well as for the study of the sensitivity of these air showers to any of the current concerns about the highest energy cosmic rays such as the expected acceptance of Pierre Auger Observatories \\cite{Auger}.  {\\bf Acknowledgements:} We are indebted to A.A.~Watson for having introduced us into the real experimental problems of HAS and an invitation to Leeds when most of this work was initially conceived. We also thank J.A.~Hinton, G.~Parente, and A.A.~Watson for many discussions and also for comments after reading our manuscript. This work was partly supported by a joint grant from the British Council and the spanish Ministry of Education (HB1997-0175), by Xunta de Galicia (XUGA-20604A98) and by CICYT (AEN99-0589-C02-02)."
    },
    "0011/astro-ph0011173_arXiv.txt": {
        "abstract": "We present {\\it Far Ultraviolet Spectroscopic Explorer} observations of the atomic and molecular absorption in high velocity cloud HVC\\,287.5+22.5+240, which lies in front of the ultraviolet-bright nucleus of the Seyfert 1 galaxy NGC\\,3783. We detect H$_2$, \\ion{N}{1}, \\ion{N}{2}, \\ion{Si}{2}, and \\ion{Fe}{2} and set limits on the amount of absorption due to \\ion{P}{3}, \\ion{Ar}{1}, and \\ion{Fe}{3}. We extend the earlier metallicity and dust-depletion measurements made by Lu and collaborators by examining the relative gas-phase abundances of Si, P, S, and Fe.   Corrections to the derived gas-phase abundances due to ionized gas in the HVC are small ($\\lesssim15$\\%). The HVC has a metallicity of 0.2--0.4 solar, similar to that of the Small Magellanic Cloud. The relative abundance pattern for the elements studied resembles that of warm gas in the SMC, which supports the idea that this HVC is part of the tidally stripped Leading Arm of the Magellanic Stream. The abundance pattern implies that the HVC contains dust grains that have been processed significantly; it is likely that the grain mantles have been modified or stripped back to expose the grain cores. We have identified more than 30 lines of H$_2$ arising in the HVC from rotational levels J\\,=\\,0 to J\\,=\\,3. Synthetic spectra and a curve-of-growth fit to these lines with b = 12 \\kms\\ indicate that log\\,N(H$_2$) = 16.80$\\pm0.10$ and f$_{H_2}$\\,=\\,2N(H$_2$)/[N(\\ion{H}{1})+2N(H$_2$)] = $1.6\\times10^{-3}$. A two-component temperature distribution is necessary to explain the observed populations of the H$_2$ rotational levels.  We find $T_{01} = 133\\pm^{37}_{21}$\\,K, and $T_{23} = 241\\pm^{20}_{17}$\\,K, indicating that the conditions in the molecular gas are more similar to those found for diffuse molecular clouds in the Galactic halo than to those for molecular clouds in the Galactic disk. From an analysis of the $J = 2$ and $J = 3$ populations, we find an absorption rate (at 1000\\,\\AA) of $\\beta_{\\rm uv} <0.1$ times the average value in the solar neighborhood. The presence of molecular gas in the HVC  requires that either the H$_2$ formed in situ or that molecules formed within the SMC survived tidal stripping.  We favor the latter possibility because of the long H$_2$ formation time ($\\sim10^8$ years) derived for this HVC. ",
        "introduction": "The spatial distribution and velocities of neutral interstellar clouds moving at high velocities ($|V_{LSR}| > 100$ \\kms) are well characterized down to \\ion{H}{1} column density levels of $\\sim10^{18}$ cm$^{-2}$, but the fundamental properties of most high velocity clouds (HVCs) remain highly uncertain because of limited information about their distances, physical properties, elemental abundances, and dust content (see Wakker \\& van~Woerden 1997 for a review).  However, for HVC Complex~C and the Magellanic Stream, ultraviolet absorption-line spectroscopy has yielded important information about the elemental composition of the gas and has shed new light on the origins of these HVCs. HVC Complex C lies at a distance d $\\gtrsim $10 kpc, has a metallicity of $\\sim$0.1--0.3 solar, and apparently contains little dust (Wakker et al. 1999; Gibson et al. 2000b; Murphy et al. 2000), which indicates that it is not material ejected into the halo as a consequence of supernovae in the Galactic disk.  Rather, it is probably metal-poor infalling gas that may be mixing with circulating Galactic gas at large distances from the Galactic plane. In contrast, \\ion{H}{1} HVC\\,287.5+22.5+240 has a metallicity (S/H = $0.25\\pm0.07$ solar, Lu et al. 1994a), dust content (Fe/H $\\approx$ 0.033 solar, Lu et al. 1998), and velocity that suggest it is associated with the Magellanic Stream, an extended ($\\sim$100\\degr) trail of gas pulled out of the Magellanic Clouds during a recent close encounter with the Galaxy (Wannier \\& Wrixon 1972; Mathewson, Cleary, \\& Murray 1974).  In this article we report on high-resolution ($\\lambda/\\Delta\\lambda \\approx 20,000$) far-ultraviolet (far-UV) absorption-line data obtained for HVC\\,287.5+22.5+240 with the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE), a powerful new observatory for studying the molecular, atomic, and ionized gases in HVCs and other interstellar/intergalactic clouds (see Sembach 1999; Oegerle et al. 2000; Sembach et al. 2000b; Shull et al. 2000a). A map of the high velocity \\ion{H}{1} in the general direction of NGC\\,3783 is shown in Figure~1. A fortuitous alignment of HVC\\,287.5+22.5+240 and the nucleus of NGC\\,3783 ($l = 287.46\\degr$, $b = +22.95\\degr$), a Seyfert 1 galaxy with a bright ultraviolet continuum, makes it possible to study the HVC in absorption below 1200\\,\\AA\\ with FUSE.  The HVC has been studied previously in \\ion{H}{1} 21\\,cm radio emission (Mathewson et al. 1974; Hulsbosch 1975; Morras \\& Bajaja 1983; Wakker et al. 1999) and at longer UV wavelengths with the {\\it Hubble Space Telescope} (HST) by Lu et al. (1994a, 1998)\\footnotemark. \\footnotetext{This HVC is identified as HVC\\,187 in the Wakker \\& van~Woerden (1991) catalog of high velocity clouds.} The HVC position and velocity of +240 \\kms\\ are consistent with gas arising in the Leading Arm of the Magellanic Stream.  The Leading Arm is predicted by N-body simulations tracing the tidal interactions of the Clouds and the Galaxy (Gardiner \\& Noguchi 1996) and has been identified in \\ion{H}{1} 21\\,cm emission (Putman et al. 1998; Putman 2000).  The FUSE spectrum of NGC\\,3783 contains unique signatures of the absorption within HVC 287.5+22.5+240, including the first detections of H$_2$ in the Magellanic Stream.  The spectrum also allows for a quantitative estimate of the ionized gas content of the HVC to be made, which is crucial for determining accurate abundances. This is particularly important in light of recent FUSE results showing that many HVCs have a collisionally ionized component seen in \\ion{O}{6} absorption.  This highly ionized gas probably occurs at either the interfaces or mixing layers between neutral HVC gas and a hot Galactic halo or within cooling regions of hot gas clouds as they are accreted onto the Galaxy (Sembach et al. 2000b).  This paper is organized as follows:  Section 2 contains a summary of the FUSE observations and data processing.  In \\S\\,3 we describe the HVC absorption and our measurements of various species.  In \\S\\,4 we consider the atomic and molecular components of the HVC and provide a short description of the intermediate velocity gas along the sight line. In \\S\\,5 we discuss the FUSE results and their implications for the origin and evolution of HVC\\,287.5+22.5+240 and the Magellanic Stream. ",
        "conclusions": "\\subsection{The Origin and Properties of HVC\\,287.5+22.5+240}  The elemental abundances, kinematics, and position of HVC\\,287.5+22.5+240 all indicate that the HVC is part of the Leading Arm of the Magellanic Stream.  The tidal plus weak drag model proposed by Gardiner (1999) shows that the material in the Leading Arm was stripped primarily from the disk of the SMC and is located a distance of about 50--75 kpc from the Sun. This feature is observed in sensitive \\ion{H}{1} 21\\,cm maps of the Magellanic Stream (Putman et al. 1998, 2000).  A detailed study of the high velocity \\ion{H}{1} 21\\,cm emission in the direction of NGC\\,3783 at different angular resolutions by Wakker et al. (1999) reveals that discrete structures exist on scales down to 1\\arcmin\\ (= 14.5 pc at a distance  D = 50 kpc).  They estimate that the HVC is part of an ensemble of small gas clouds with internal thermal pressures, $P/k \\sim 18,000$\\,D$^{-1}_{kpc}$\\,K\\,cm$^{-3}$, implying $P/k = 1.1\\,n_H\\,T \\sim360$ K cm$^{-3}$ for an assumed distance of 50 kpc.  The number of high velocity clouds for which both gas-phase abundances and molecular content are available is very limited.  The two highest quality measurements to date are for HVC\\,287.5+22.5+240 (the subject of this paper) and Complex~C.  For HVC\\,287.5+22.5+240, we find N(H$_2$) = $6.3\\times10^{16}$ cm$^{-2}$ and f$_{H_2}$ = $1.6\\times10^{-3}$.  We infer the presence of dust from the relative abundances of Si, S, and Fe (\\S\\,4.1). For Complex~C in the direction of Mrk~876 ($l=98.27\\degr, b = +40.38\\degr$), Murphy et al. (2000) find N(H$_2$) $<2\\times10^{14}$ cm$^{-2}$ and f$_{H_2}$ $< 1.7\\times10^{-5}$.  Combined with the abundance information available for Fe from FUSE (Murphy et al. 2000) and sulfur from HST (Wakker et al. 1999; Gibson et al. 2000b), it appears that Complex~C has a low metallicity ($\\sim0.1-0.3$ solar) and essentially no depletion of Fe onto dust grains.  Thus, these two cases imply that dust and H$_2$ are related in HVCs -- low values of f$_{H_2}$ and mild depletions appear to go together.  High velocity H$_2$ has also been reported for HVCs in the direction of the LMC, although the detections are not as secure as the detection for HVC\\,287.5+22.5+240. Using the echelle spectrograph on the {\\it Orbiting and Retrievable Far and Extreme Ultraviolet Spectrometers} (ORFEUS), Richter et al. (1999) found N(H$_2$) $=(2.2-3.6)\\times10^{15}$ cm$^{-2}$ and f$_{H_2}$ $=(3.6-6.0)\\times10^{-4}$ in the +120 \\kms\\ HVC toward the LMC star HD\\,269546 ($l = 279.32\\degr, b = -32.77\\degr$)\\footnote{As this paper was being refereed, HD\\,269546 (Sk\\,-68\\,82) was observed with FUSE.  The FUSE spectra indicate that the H$_2$ absorption is weaker than deduced from the ORFEUS data (P. Richter, private communication).  This suggests that f$_{H_2}$ for the +120 \\kms\\ gas in the LMC direction may be less than the value stated above.}. Their analysis of {\\it International Ultraviolet Explorer} (IUE) \\ion{Fe}{2} data indicated that (Fe/H) in the cloud is roughly half the solar value. Several higher quality observations of other LMC stars by FUSE also reveal H$_2$ absorption and will allow a more thorough investigation of the composition and molecular fraction of the high velocity gas in front of the LMC (Richter et al. 2001).  \\subsection{Where did the H$_2$ Originate?}  The presence of molecular gas in HVC\\,287.5+22.5+240 requires that either the H$_2$ formed in situ or that the H$_2$ was originally formed in the SMC and survived the tidal stripping that created the Leading Arm.    For the six LMC and one SMC sight lines for which Magellanic Cloud H$_2$ was observed with ORFEUS or the {\\it Hopkins Ultraviolet Telescope} (HUT), four have values of f$_{H_2}$ $< 1\\times10^{-4}$, one has a value of $6.6\\times10^{-3}$, and two have values $> 3.8\\times10^{-2}$  (see Richter 2000 and references therein). FUSE observations of a much larger sample of stars ($\\sim40$) demonstrate that  H$_2$ is often present in the Clouds; preliminary results for 10 stars in both the LMC and SMC (Tumlinson et al. 2001) yield $\\langle$f$_{H_2}$$\\rangle$ $\\approx$ $1.5\\times10^{-2}$ in the LMC and $\\langle$f$_{H_2}$$\\rangle$ $\\approx$ $9.0\\times10^{-3}$ in the SMC, with a range of values encompassing both the low and high f$_{H_2}$ branches. However, few points fall in the ``transition'' region between low and high f$_{H_2}$.  In this respect, the Magellanic Cloud molecular hydrogen fractions resemble those found for Milky Way disk and halo gas (Savage et al. 1977; Shull et al. 2000b). If the H$_2$ in HVC\\,287.5+22.5+240 is remnant material stripped from the SMC, it likely originated in a region having a higher value of f$_{H_2}$ than the presently observed value of $1.6\\times10^{-3}$.  The formation and destruction of H$_2$ are influenced by many competing factors.  These include the quantity and type of dust grains, the gas density, the amount of self-shielding by the H$_2$, the intensity of the far-ultraviolet radiation field, the presence of hot gas, and the effects of shocks. At the column density levels observed, it is unlikely that the H$_2$ in the HVC is thermalized, and in a dynamic environment such as the Magellanic Stream the competing production and destruction effects may not be equilibrated. Thus, it is difficult to quantify the net H$_2$ production rate.  However, it is still useful to briefly consider what some of the competing effects may be.  Dust grains serve as catalysts for the formation of H$_2$ (Hollenbach, Werner, \\& Salpeter 1971; Spitzer \\& Jenkins 1975) and also act as a shielding agent against ultraviolet photons capable of dissociating the molecules. The formation of H$_2$ proceeds most readily for high grain surface area and cold gas and grain temperatures. The heavy element abundance pattern in the HVC is similar to that in warm (T $\\sim 10^3$ K) clouds located in the Galactic halo, which suggests that the grains have been stripped of their outer mantles and consist mainly of their residual (oxide) cores (Sembach \\& Savage 1996). Thus, much of the grain surface has probably been modified and may be less conducive to H$_2$ formation than the surfaces of grains found in cold molecular clouds in the Galactic disk.  The ultraviolet radiation field at the location of the Leading Arm is probably less intense than for typical interstellar environments within either the Galaxy or the SMC.  The exact details of the radiation field and heating of the gas depend upon the transfer of photons out of the Galaxy and the Magellanic Clouds (see Wolfire et al. 1995; Bland-Hawthorn \\& Maloney 1999).      The rotational excitation temperatures derived for HVC\\,287.5+22.5+240 are consistent with modest photon pumping of the $J = 2$ and 3 levels but are not indicative of very high levels of excitation like those found for clouds near hot stars (no higher rotational level lines are seen).  To date, no early-type stars have been found in the Magellanic Stream. In the Galaxy, much higher excitation temperatures ($T_{ex} > 500$ K) are often found for H$_2$ in the vicinity of hot stars, and similarly high values have been observed for interstellar gas in the Magellanic Clouds (e.g., Friedman et al. 2000; Mallouris et al. 2001; Tumlinson et al. 2001).  The absorbing clouds have 21\\,cm emission structure on the scale of $\\sim 1\\arcmin - 10\\arcmin$ (Putman \\& Gibson 1999; Wakker et al. 1999), which corresponds to sizes of $\\sim 1-100$ pc at the estimated 50 kpc distance of the HVC. The corresponding temperature range is $\\sim 100 - 1000$\\,K, if the pressure model of \\S5.1 is correct.  Directly along the NGC\\,3783 sight line, the \\ion{H}{1} 21\\,cm emission has FWHM $\\approx$ 22 \\kms\\ (see Figure~2 in Lu et al. 1998).  For a line broadened solely by thermal gas motions, this corresponds to $T = 21.7 (FWHM)^2 \\approx 10^4$ K.  However, this is an upper limit since the line may also be broadened significantly by turbulence and sub-component structure.  This appears to be the case since the lines of species heavier than \\ion{H}{1} (e.g., \\ion{Si}{2}, \\ion{Fe}{2}) are also broad and have effective b-values much larger than those expected for thermal motions alone (see Figure~3 and \\S\\,4).  If we adopt a gas pressure, $P/k = 360~{\\rm cm}^{-3}$~K, and assume $T = (100~{\\rm K})T_{100}$ with He/H = 0.1, the hydrogen density is $n_H = (3.3~{\\rm cm}^{-3})(T_{100})^{-1}$, from which we estimate a characteristic cloud dimension of $L = N({\\rm H~I})/n_H = (8~{\\rm pc})(T_{100})$. The observations are therefore consistent with $0.3 < n_H < 3$ cm$^{-3}$ for $100 < T < 1000$~K. This range of densities is in the low-$n_H$ limit for models of H$_2$ radiative excitation and collisional de-excitation. Preliminary estimates based on fitting the populations of rotational levels $J = 2$ and $J=3$ suggest that the H$_2$ absorption rate in the Lyman and Werner bands (1000 \\AA) is $\\beta_{\\rm uv} < 5 \\times 10^{-11}$ s$^{-1}$, or $<0.1$ times the average value, $\\beta_0 = 5 \\times 10^{-10}$ s$^{-1}$ in the solar neighborhood (Jura 1974).  The other general statement that can made concerning the origin of H$_2$ in this HVC is the characteristic formation time. The H$_2$ formation rate depends on the collision rate between \\ion{H}{1} atoms and grains, the probability that \\ion{H}{1} atoms are adsorbed on the grain surface (``sticking probability''), the mobility and lifetime of the atoms on the surface, and the probability that a molecule is ejected after formation. The volume formation rate of H$_2$ on grain surfaces (Hollenbach \\& Salpeter 1971; Shull \\& Beckwith 1982) is written as $R n_H$, where $R \\approx (1-3) \\times 10^{-17}$ cm$^3$~s$^{-1}$ in the local ISM.  Molecule formation proceeds rapidly for grains with surface temperatures less than a critical value, $T_{\\rm cr} \\approx 20-100$ K;  H$_2$ formation may also be suppressed when $T_{\\rm HI} \\gg 100$~K. For the range of physical conditions for this HVC, we estimate the H$_2$ formation time to be, \\begin{equation} t_{\\rm form} \\approx \\left[ R n_{\\rm HI} \\right] ^{-1} \\approx  1 \\times 10^8 ~{\\rm yr} \\end{equation} for $R \\approx 10^{-17}$ cm$^3$~s$^{-1}$ and $n_H = n_{\\rm HI} \\approx 3$ cm$^{-3}$. Although this time is a significant fraction of the ($\\sim1$ Gyr) orbital period of the Magellanic Stream, new H$_2$ could form up to the observed fraction, $f_{\\rm H2} \\approx 10^{-3}$, on the surfaces of dust grains that survived the process of tidal stripping. More likely, much of the observed H$_2$ survived stripping and has resisted photodestruction by the high-latitude ultraviolet radiation field because of self-shielding in the dissociating Lyman bands.  Self-shielding becomes important and reduces the destruction rate (\\,$\\langle k \\rangle \\beta_{\\rm uv}$\\,) by a substantial factor once N(H$_2$) $\\gg$ 10$^{14}$ cm$^{-2}$ (optical depth unity in the Lyman band).  \\subsection{Relevance to QSO Absorption Lines} Absorption line spectroscopy of the Lyman and Werner bands of H$_2$ provides the most sensitive means available for assessing the molecular content of cold gas in the Universe. To date, detections of H$_2$ absorption outside the Milky Way have been limited to the Magellanic Clouds (see Richter 2000; Shull et al. 2000b; Mallouris et al. 2001; Tumlinson et al. 2001), a few high redshift ($z \\gtrsim 2$) quasar absorption line systems (Levshakov \\& Varshalovich 1985; Ge \\& Bechtold 1997; Levshakov et al. 2000), and now HVC\\,287+22.5+240.  Although a few low column density quasar absorption line systems have modest values of f$_{H_2}$, most have very little molecular gas and do not appear to sustain a large quantity of molecules due to the intense radiation field impinging on the clouds (see Black, Chaffe, \\& Foltz 1987; Levshakov et al. 2000). These systems also contain little dust. HVCs represent the best low-redshift analogs to the QSO absorption line systems with 17 $\\lesssim$ log(\\ion{H}{1}) $\\lesssim$ 19.  Therefore, a better understanding of the processes that affect the elemental abundances, molecular content, and dust composition of HVCs could provide important insights into similar quantities derived for the quasar absorption line systems.  This is especially important for HVCs that originate either from tidal interactions (as appears to be the case for HVC\\,287.5+22.5+240) or are part of a larger ensemble of extragalactic clouds in the Local Group that may be interacting with the Milky Way (see, e.g., Blitz et al. 1999; Sembach et al. 1999, 2000b). If such dust-bearing clouds commonly exist around other galaxies, particularly those at high redshift, the dust could produce intragroup reddening that might affect studies of objects in those distant galaxies (e.g., the dimming of very distant supernovae used to study the expansion of the Universe).    \\smallskip We thank the FUSE Science and Operations Teams for their dedicated efforts to make the observations described in this paper possible. We also thank Mary Putman and Brad Gibson for providing a preliminary electronic version of Figure 1, Philipp Richter and Daniel Welty for useful comments on the manuscript, and  Matthew Browning for assistance with the calculation of $\\beta_{\\rm uv}$. This work is based on data obtained for the FUSE Science Team by the NASA-CNES-CSA FUSE mission operated by the Johns Hopkins University. Financial support has been provided by NASA contract NAS5-32985.  KRS and JCH acknowledge partial support from NASA Long Term Space Astrophysics grant NAG5-3485.  JMS acknowledges support from NASA LTSA grant NAG5-7262.  \\clearpage \\newpage"
    }
}